EVLA Guides evlaguides_nrao_edu https://evlaguides.nrao.edu/index.php?title=Main_Page MediaWiki 1.38.6 first-letter Media Special Talk User User talk EVLA Guides EVLA Guides talk File File talk MediaWiki MediaWiki talk Template Template talk Help Help talk Category Category talk 0 1 1 2010-04-07T16:26:15Z MediaWiki default 0 wikitext text/x-wiki <big>'''MediaWiki has been successfully installed.'''</big> Consult the [http://meta.wikimedia.org/wiki/Help:Contents User's Guide] for information on using the wiki software. == Getting started == * [http://www.mediawiki.org/wiki/Manual:Configuration_settings Configuration settings list] * [http://www.mediawiki.org/wiki/Manual:FAQ MediaWiki FAQ] * [https://lists.wikimedia.org/mailman/listinfo/mediawiki-announce MediaWiki release mailing list] bd962048d95fbb6b6b514885867811db20a5476b 2 1 2010-04-08T14:24:00Z Jmcmulli 2 wikitext text/x-wiki {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:center" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 2: What's New --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:center" | <div style="font-size:120%">What's New</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{What's New}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| {{Latest News}} {{Getting Started}} {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:center;" | <div style="font-size:120%">Collaborate</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| |- | style="padding-left:6px; padding-top:6px; padding-bottom:2px; padding-right:0px; font-size:8pt" | * '''Create a new article''': See above. * [[Special:Wantedpages|Wanted Pages]] * [[:Category:Pages needing further work|Pages needing further work]] * [[:Category:Pages needing expert attention|Pages needing expert attention]] * [http://www.mediawiki.org/wiki/Manual:Configuration_settings Configuration settings list] * [http://www.mediawiki.org/wiki/Manual:FAQ MediaWiki FAQ] * [https://lists.wikimedia.org/mailman/listinfo/mediawiki-announce MediaWiki release mailing list] * Consult the [http://meta.wikimedia.org/wiki/Help:Contents User's Guide] for information on using the wiki software. |} <!-- TEXT 2 (END) --> |} |} Welcome to EVLA Guides wiki page. 501f1e874cdba8c1bd7877b0bc9a36bbd72f2584 37 2 2010-04-13T15:02:42Z Jmcmulli 2 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides</div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:center" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 2: What's New --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:center" | <div style="font-size:120%">What's New</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{What's New}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| {{Latest News}} {{Getting Started}} {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:center;" | <div style="font-size:120%">Collaborate</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| |- | style="padding-left:6px; padding-top:6px; padding-bottom:2px; padding-right:0px; font-size:8pt" | * '''Create a new article''': See above. * [[Special:Wantedpages|Wanted Pages]] * [[:Category:Pages needing further work|Pages needing further work]] * [[:Category:Pages needing expert attention|Pages needing expert attention]] * [http://www.mediawiki.org/wiki/Manual:Configuration_settings Configuration settings list] * [http://www.mediawiki.org/wiki/Manual:FAQ MediaWiki FAQ] * [https://lists.wikimedia.org/mailman/listinfo/mediawiki-announce MediaWiki release mailing list] * Consult the [http://meta.wikimedia.org/wiki/Help:Contents User's Guide] for information on using the wiki software. |} <!-- TEXT 2 (END) --> |} |} Welcome to EVLA Guides wiki page. e3235cf7a384bc960bbf1b545988bc940bba2781 39 37 2010-04-13T15:07:22Z Jmcmulli 2 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 2: What's New --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">What's New</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{What's New}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| {{Latest News}} {{Getting Started}} {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:left;" | <div style="font-size:120%">Collaborate</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| |- | style="padding-left:6px; padding-top:6px; padding-bottom:2px; padding-right:0px; font-size:8pt" | * '''Create a new article''': See above. * [[Special:Wantedpages|Wanted Pages]] * [[:Category:Pages needing further work|Pages needing further work]] * [[:Category:Pages needing expert attention|Pages needing expert attention]] * [http://www.mediawiki.org/wiki/Manual:Configuration_settings Configuration settings list] * [http://www.mediawiki.org/wiki/Manual:FAQ MediaWiki FAQ] * [https://lists.wikimedia.org/mailman/listinfo/mediawiki-announce MediaWiki release mailing list] * Consult the [http://meta.wikimedia.org/wiki/Help:Contents User's Guide] for information on using the wiki software. |} <!-- TEXT 2 (END) --> |} |} Welcome to EVLA Guides wiki page. bcbcf7f8b11a2a7761ebcc8600be96efe4ff2190 41 39 2010-04-13T15:27:23Z Jmcmulli 2 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">Featured Article</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{Featured Article}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:left;" | <div style="font-size:120%">News</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {{What's New}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:left;" | <div style="font-size:120%">Configuration & Proposal Dates</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {{What's New}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> {{Latest News}} {{EVLA Configuration & Proposal Timelines}} 5650cd591f407df3747b017a3708dc8e9d558afe 44 41 2010-04-13T15:38:30Z Jmcmulli 2 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">Featured Article</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{Featured Article}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:left;" | <div style="font-size:120%">News</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> {{News}} <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:left;" | <div style="font-size:120%">Configuration & Proposal Dates</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> {{Configuration & Proposal Timelines}} 41687384a6b605cb6c131a76d0df7bcd4ec52578 47 44 2010-04-13T17:08:06Z Jmcmulli 2 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{Explore ALMA}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">Featured Article</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{Featured Article}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">News</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | * 12-Apr-2010: First 27-antenna correlation {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:left;" | <div style="font-size:120%">Configuration & Proposal Dates</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| |- | style="padding-left:6px; padding-top:6px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| ! Trimester ! Observing Period ! Configuration ! Proposal Deadline |- | 2010a | 2010 Mar 01 - 2010 May 24 | align="center" | D | 2009 Oct 1 |- | 2010a | 2010 May 28 - 2010 Jun 14 | align="center" | DnC | 2009 Oct 1 |- | 2010b | 2010 Jun 25 - 2010 Sep 13 | align="center" | C | 2010 Feb 1 |- | 2010b | 2010 Sep 17 - 2010 Oct 04 | align="center" | CnB | 2010 Feb 1 |- | 2010c | 2010 Oct 15 - 2011 Jan 03 | align="center" | B | 2010 Jun 1 |- | 2010c | 2011 Jan 07 - 2011 Jan 24 | align="center" | BnA | 2010 Jun 1 |- |} <!-- TEXT 2 (END) --> |} |} a08e104c2552c191bc6e9096dbac9d682b8c465c 48 47 2010-04-13T17:12:01Z Jmcmulli 2 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">Featured Article</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{Featured Article}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">News</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | * 12-Apr-2010: First 27-antenna correlation {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:left;" | <div style="font-size:120%">Configuration & Proposal Dates</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| |- | style="padding-left:6px; padding-top:6px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| ! Trimester ! Observing Period ! Configuration ! Proposal Deadline |- | 2010a | 2010 Mar 01 - 2010 May 24 | align="center" | D | 2009 Oct 1 |- | 2010a | 2010 May 28 - 2010 Jun 14 | align="center" | DnC | 2009 Oct 1 |- | 2010b | 2010 Jun 25 - 2010 Sep 13 | align="center" | C | 2010 Feb 1 |- | 2010b | 2010 Sep 17 - 2010 Oct 04 | align="center" | CnB | 2010 Feb 1 |- | 2010c | 2010 Oct 15 - 2011 Jan 03 | align="center" | B | 2010 Jun 1 |- | 2010c | 2011 Jan 07 - 2011 Jan 24 | align="center" | BnA | 2010 Jun 1 |- |} <!-- TEXT 2 (END) --> |} |} 5c2209ea79e60795870932db59b62af6febb6393 10 2 3 2010-04-08T14:35:30Z Jmcmulli 2 Created page with '{|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |[http://science.nrao.edu/evla/ '''EVLA Information'''] · [https://staff.nrao.edu/wiki/bin/view/EVL…' wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |[http://science.nrao.edu/evla/ '''EVLA Information'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification'''] |- |valign=top|[[Image:book.gif]] '''[[:Category:Technical Documentation|Technical Documentation]]'''<br>[[Antennas]] · [[Receivers]] |- |valign=top|[[Image:technical.gif]] '''[[:Category:Calibration|Calibration Data]]'''<br> [[:Category:Calibration|Flux]] · [[:Category:Calibration|Polarization]] |- |valign=top|[[Image:technical.gif]] '''[[:Category:Post-Processing|Post-Processing]]'''<br> [[:Category:Post-Processing|Reduction Strategies]] 02dba7a85b5612320200683ea849ad949b2b45b3 6 3 2010-04-08T14:53:03Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 [[File:vla_panorama_lo.jpg|300px|left]] |- |[http://science.nrao.edu/evla/ '''EVLA Information'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification'''] |- |valign=top|[[Image:book.gif]] '''[[:Category:Technical Documentation|Technical Documentation]]'''<br>[[Antennas]] · [[Receivers]] |- |valign=top|[[Image:technical.gif]] '''[[:Category:Calibration|Calibration Data]]'''<br> [[:Category:Calibration|Flux]] · [[:Category:Calibration|Polarization]] |- |valign=top|[[Image:technical.gif]] '''[[:Category:Post-Processing|Post-Processing]]'''<br> [[:Category:Post-Processing|Reduction Strategies]] 8d7f3fb82218f407b8162f4e8e79147d4bd43a85 8 6 2010-04-08T14:57:04Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 [[File:vla_panorama_lo.jpg|300px|center]] |- |[http://science.nrao.edu/evla/ '''EVLA Information'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification'''] |- |valign=top|[[Image:book.gif]] '''[[:Category:Technical Documentation|Technical Documentation]]'''<br>[[Antennas]] · [[Receivers]] |- |valign=top|[[Image:technical.gif]] '''[[:Category:Calibration|Calibration Data]]'''<br> [[:Category:Calibration|Flux]] · [[:Category:Calibration|Polarization]] |- |valign=top|[[Image:technical.gif]] '''[[:Category:Post-Processing|Post-Processing]]'''<br> [[:Category:Post-Processing|Reduction Strategies]] 88e207a858a783da5ecce0c908d3e7fc9c6f8069 9 8 2010-04-12T16:39:49Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 [[File:vla_panorama_lo.jpg|300px|center]] |- |[http://science.nrao.edu/evla/ '''EVLA Information'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification'''] |- |valign=top|[[Image:book.gif]] '''[[:Category:Post-Processing|Post-Processing]]'''<br> [[:Category:Post-Processing|Reduction Strategies]] |- |valign=top|[[Image:book.gif]] '''[[:Category:Calibration|Calibration Data]]'''<br> [[:Category:Calibration|Flux]] · [[:Category:Calibration|Polarization]] |- |valign=top|[[Image:book.gif]] '''[[:Category:Technical Documentation|Technical Documentation]]'''<br>[[Antennas]] · [[Receivers]] 9382d04ef824355ab8adb931ae58ddd67b1b855b 10 9 2010-04-12T16:45:24Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 [[File:vla_panorama_lo.jpg|300px|center]] |- |[http://science.nrao.edu/evla/ '''EVLA Information'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification'''] |- |valign=top|[[Image:book.gif]] '''[[:Category:Post-Processing|Post-Processing]]'''<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] |- |valign=top|[[Image:book.gif]] '''[[:Category:Calibration|Calibration Data]]'''<br> [[:Category:Calibration|Flux]] · [[:Category:Calibration|Polarization]] |- |valign=top|[[Image:book.gif]] '''[[:Category:Technical Documentation|Technical Documentation]]'''<br>[[Antennas]] · [[Receivers]] 679e5759b24ce009b8b8489a7bb9f7cf1ba50d5f 38 10 2010-04-13T15:04:24Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |[http://science.nrao.edu/evla/ '''EVLA Information'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification'''] |- |valign=top|[[Image:book.gif]] '''[[:Category:Post-Processing|Post-Processing]]'''<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] |- |valign=top|[[Image:book.gif]] '''[[:Category:Calibration|Calibration Data]]'''<br> [[:Category:Calibration|Flux]] · [[:Category:Calibration|Polarization]] |- |valign=top|[[Image:book.gif]] '''[[:Category:Technical Documentation|Technical Documentation]]'''<br>[[Antennas]] · [[Receivers]] 64f938bbdd65a5da4bfeda0692fa35ed1720e2c9 40 38 2010-04-13T15:14:04Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] a837ec208f2f4dfdde9e1904d0351083decacf3c 45 40 2010-04-13T15:42:38Z Jmcmulli 2 wikitext text/x-wiki {|width=52% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] 8685bee68da6aea69e23ffa7810e2d6810695192 46 45 2010-04-13T15:44:56Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] a837ec208f2f4dfdde9e1904d0351083decacf3c 6 3 4 2010-04-08T14:42:42Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 5 4 2010-04-08T14:44:59Z Jmcmulli 2 uploaded a new version of "[[File:Book.gif]]" wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 4 7 2010-04-08T14:54:55Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 14 5 11 2010-04-12T17:00:50Z Jmcmulli 2 Created page with '[[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] == Special Considerations for EVLA data calibration and imaging (AIPS) == The old VLA with its once state…' wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] == Special Considerations for EVLA data calibration and imaging (AIPS) == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., d12a9c451ce2f5c2afb842b44a8e401aef9dae99 12 11 2010-04-12T19:15:36Z Jmcmulli 2 wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang="text"> > DEFAULT ’UVLOD’ ; INP CR > DATAIN ’E2E:filename’ CR > DOUVCOMP FALSE CR > OUTNA ’myname’ CR > OUTCL ’ ’ CR > OUTSEQ 0 CR > OUTDI 3 CR > INP CR > GO CR </source> to initialize and review the inputs needed. where filename is the disk file name in logical area E2E; (see § 3.10.3). to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. to set the AI P S name. to take default (UVDATA) class. to take next higher sequence #. to write the data to disk 3 (one with enough space). to review the inputs. to run the program when you’re satisfied with inputs. 56cde69eb687ecde43371b313deea43e90ca6d44 13 12 2010-04-12T19:25:34Z Jmcmulli 2 wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <pre> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </pre> 5cba02a114fee20bfd0aceb3ffd4f2afd5b840f1 14 13 2010-04-12T19:28:11Z Jmcmulli 2 wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <pre> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </pre> == Initial Calibration FRING == fd10b04a827fb076ff254699fe4bfd156e1bcd97 15 14 2010-04-12T22:50:54Z Jmcmulli 2 wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> == Initial Calibration FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. > TASK ’CLCAL’ ; INP CR > TIMERANG 0 CR > GAINUSE 0 ; GAINVER 0 CR to look at the necessary inputs. to reset the time range. to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). 2d7de44e82364773c94d2882a4a4622af3c22e8c 16 15 2010-04-12T23:29:41Z Jmcmulli 2 wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> </source> == Basic Calibration == == Target Source Data -- Edit and SPLIT == == Spectral-line Imaging Hints == == Continuum Imaging Hints == == Concluding Remakrs, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] 68d546851f7d078e83ab6fa699a18c0662d82d10 18 16 2010-04-12T23:43:38Z Jmcmulli 2 wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == == Target Source Data -- Edit and SPLIT == == Spectral-line Imaging Hints == == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] 17c155cd540fe084b70c0a8613b5ee20d3987dc0 19 18 2010-04-13T13:11:16Z Jmcmulli 2 /* Initial Editing */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == == Target Source Data -- Edit and SPLIT == == Spectral-line Imaging Hints == == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] c59cf0e26aed22deead9b39134ae8bb47e262f38 20 19 2010-04-13T13:17:24Z Jmcmulli 2 /* Basic Calibration */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP > INDI n; GETN m CR > DOCAL 1 CR > SOLINT 0 CR > CALSOUR ’bandpass cal’ CR to reset all adverbs and choose the task. to select the data set on disk n and catalog number m. to apply the delay calibration — very important. to compute a bandpass solution for each scan on the bandpass calibrator. to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR > GO CR </source> averaging in each IF. Remember these values — you will use them again. to normalize the results only after the solution is found using the channels selected by ICHANSEL. to make a bandpass (BP) table. Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: > DEFAULT CALIB ; INP > INDI n; GETN m CR > IN2DI n2; GET2N m2 CR to reset all adverbs and choose the task. to select the data set on disk n and catalog number m. to select the model image on disk n2 and catalog number m2. E – 6 E. Handling EVLA Data in AIPS >DOCAL1;DOBAND3CR > SOLINT 0 ; NMAPS 1 CR > CALSOUR ’flux cal’ CR AIPS CookBook: 31-December-2010 (revised 19-January-2010) E.5. Target source data — edit and SPLIT to apply the delay and bandpass calibration — very important. to compute a solution for each calibration scan and use the source model. to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: > CLR2NAME ; NMAPS 0 CR > CALSOUR ’other cal1’, ’other cal2’ CR > GO CR to do no models. to select the secondary calibrators by whatever names appear in your LISTR output. to find the remaining complex gains. Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: > TASK ’GETJY’ ; INP CR > SOURCE CALSOUR CR > CALSOUR ’flux cal’ CR > INP CR > GO CR to set the task name without changing other adverbs. to select the secondary sources by the list of name you just used. to select the primary flux calibrator by whatever form of its name appears in your LISTR output. to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. to adjust the gains in the SN table and the fluxes in the SU (source) table. Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: > DEFAULT CLCAL ; INP CR > INDI n; GETN m CR > CALCODE ’*’ CR > SNVER 2; INVERS SNVER CR > GO CR Check the result using POSSM and/or VPLOT. to clear the adverbs. to select the data set on disk n and catalog number m. to select all calibration sources. to select your solution table from CALIB. Do not include the SN table from FRING a second time! to apply SN table 2 to CL table 2, creating CL table 3. == Target Source Data -- Edit and SPLIT == == Spectral-line Imaging Hints == == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] 8872155793ec7129466c869c6b23acf39fee22c4 21 20 2010-04-13T13:48:00Z Jmcmulli 2 /* Basic Calibration */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 CR to apply the delay calibration — very important. > SOLINT 0 CR to compute a bandpass solution for each scan on the bandpass calibrator. > CALSOUR ’bandpass cal’ CR to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. Remember these values — you will use them again. > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR to normalize the results only after the solution is found using the channels selected by ICHANSEL. > GO CR to make a bandpass (BP) table. </source> Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: <source lang="text"> > DEFAULT CALIB ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IN2DI n2; GET2N m2 CR to select the model image on disk n2 and catalog number m2. > DOCAL1;DOBAND3 CR to apply the delay and bandpass calibration — very important. > SOLINT 0 ; NMAPS 1 CR to compute a solution for each calibration scan and use the source model. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. </source> Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: <source lang="text"> > CLR2NAME ; NMAPS 0 CR to do no models. > CALSOUR ’other cal1’, ’other cal2’ CR to select the secondary calibrators by whatever names appear in your LISTR output. > GO CR to find the remaining complex gains. </source> Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: <source lang="text"> > TASK ’GETJY’ ; INP CR to set the task name without changing other adverbs. > SOURCE CALSOUR CR to select the secondary sources by the list of name you just used. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > INP CR to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. > GO CR to adjust the gains in the SN table and the fluxes in the SU (source) table. </source> Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: <source lang="text"> > DEFAULT CLCAL ; INP CR to clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > CALCODE ’*’ CR to select all calibration sources. > SNVER 2; INVERS SNVER CR to select your solution table from CALIB. Do not include the SN table from FRING a second time! > GO CR to apply SN table 2 to CL table 2, creating CL table 3. </source> Check the result using POSSM and/or VPLOT. == Target Source Data -- Edit and SPLIT == == Spectral-line Imaging Hints == == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] d228f1c2473e60f650505319ad7449a6f763209e 22 21 2010-04-13T14:03:42Z Jmcmulli 2 /* Target Source Data -- Edit and SPLIT */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 CR to apply the delay calibration — very important. > SOLINT 0 CR to compute a bandpass solution for each scan on the bandpass calibrator. > CALSOUR ’bandpass cal’ CR to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. Remember these values — you will use them again. > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR to normalize the results only after the solution is found using the channels selected by ICHANSEL. > GO CR to make a bandpass (BP) table. </source> Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: <source lang="text"> > DEFAULT CALIB ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IN2DI n2; GET2N m2 CR to select the model image on disk n2 and catalog number m2. > DOCAL1;DOBAND3 CR to apply the delay and bandpass calibration — very important. > SOLINT 0 ; NMAPS 1 CR to compute a solution for each calibration scan and use the source model. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. </source> Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: <source lang="text"> > CLR2NAME ; NMAPS 0 CR to do no models. > CALSOUR ’other cal1’, ’other cal2’ CR to select the secondary calibrators by whatever names appear in your LISTR output. > GO CR to find the remaining complex gains. </source> Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: <source lang="text"> > TASK ’GETJY’ ; INP CR to set the task name without changing other adverbs. > SOURCE CALSOUR CR to select the secondary sources by the list of name you just used. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > INP CR to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. > GO CR to adjust the gains in the SN table and the fluxes in the SU (source) table. </source> Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: <source lang="text"> > DEFAULT CLCAL ; INP CR to clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > CALCODE ’*’ CR to select all calibration sources. > SNVER 2; INVERS SNVER CR to select your solution table from CALIB. Do not include the SN table from FRING a second time! > GO CR to apply SN table 2 to CL table 2, creating CL table 3. </source> Check the result using POSSM and/or VPLOT. == Target Source Data -- Edit and SPLIT == At this point, your calibration should be finished. You should now do an initial editing on the target sources, much like that done above for the calibration sources: <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all channels into one number. > CALCODE ’-CAL’ CR to do just target sources now. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> Again, remember to set it to flag all channels and IFs. You may have to select sub-windows and force the averaging to one times ∆ t to edit in detail, or perhaps the default time averaging will be beneficial. In general, the DISPLAY AMP V DIFF is a powerful way to catch bad amplitudes and phases. It will catch drop outs either as bright lines for strong sources or dark grey ones for weak sources. Since EVLA data sets tend to be large and unwieldy, it is recommended that you separate the data into the separate target sources, applying the current calibration and flagging once and for all. The imaging task IMAGR can do this on the fly, but, especially for observations of spectral-line sources, this is excessively expensive. <source lang="text"> > DEFAULT SPLIT ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > CALCODE ’-CAL’ CR to do just target sources now. > GO CR to write out separate calibrated data sets for each target source. </source> At present, EVLA data sets have no meaningful weights associated with the visibility data. There is a task new to 31DEC10 called REWAY which computes a robust rms over spectral channels within each IF and polarization. It can simply base the weights on these on a record-by-record, baseline-by-baseline basis or it can average the rmses in time, solve for antenna-based rmses, smooth those further in time, and only then apply them to the data. For these weights to be meaningful, the bandpass calibration must be applied and any RFI or other real spectral-line signal channels must be omitted from the rms computation. For the weights to be correctly calibrated, all amplitude calibration must also be applied. For these reasons, REWAY might well be used instead of SPLIT, running it one source at a time. Thus, <source lang="text"> > DEFAULT REWAY ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > SOURCE ’target1’ , ’ ’ CR to do one target source. > GO CR to write out a calibrated, weighted data set for the first target source. </source> Then, when that finishes <source lang="text"> > SOURCE ’target2’ , ’ ’ ; GO CR to do another target source. </source> It is not clear that this algorithm is optimal, but it certainly should be better than using all weights 1.0 throughout. == Spectral-line Imaging Hints == == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] 48ca11675d4282343d2cfcfec1f8ed30bd0297cf 23 22 2010-04-13T14:06:19Z Jmcmulli 2 /* Target Source Data -- Edit and SPLIT */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 CR to apply the delay calibration — very important. > SOLINT 0 CR to compute a bandpass solution for each scan on the bandpass calibrator. > CALSOUR ’bandpass cal’ CR to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. Remember these values — you will use them again. > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR to normalize the results only after the solution is found using the channels selected by ICHANSEL. > GO CR to make a bandpass (BP) table. </source> Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: <source lang="text"> > DEFAULT CALIB ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IN2DI n2; GET2N m2 CR to select the model image on disk n2 and catalog number m2. > DOCAL1;DOBAND3 CR to apply the delay and bandpass calibration — very important. > SOLINT 0 ; NMAPS 1 CR to compute a solution for each calibration scan and use the source model. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. </source> Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: <source lang="text"> > CLR2NAME ; NMAPS 0 CR to do no models. > CALSOUR ’other cal1’, ’other cal2’ CR to select the secondary calibrators by whatever names appear in your LISTR output. > GO CR to find the remaining complex gains. </source> Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: <source lang="text"> > TASK ’GETJY’ ; INP CR to set the task name without changing other adverbs. > SOURCE CALSOUR CR to select the secondary sources by the list of name you just used. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > INP CR to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. > GO CR to adjust the gains in the SN table and the fluxes in the SU (source) table. </source> Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: <source lang="text"> > DEFAULT CLCAL ; INP CR to clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > CALCODE ’*’ CR to select all calibration sources. > SNVER 2; INVERS SNVER CR to select your solution table from CALIB. Do not include the SN table from FRING a second time! > GO CR to apply SN table 2 to CL table 2, creating CL table 3. </source> Check the result using POSSM and/or VPLOT. == Target Source Data -- Edit and SPLIT == At this point, your calibration should be finished. You should now do an initial editing on the target sources, much like that done above for the calibration sources: <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all channels into one number. > CALCODE ’-CAL’ CR to do just target sources now. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> Again, remember to set it to flag all channels and IFs. You may have to select sub-windows and force the averaging to one times ∆ t to edit in detail, or perhaps the default time averaging will be beneficial. In general, the DISPLAY AMP V DIFF is a powerful way to catch bad amplitudes and phases. It will catch drop outs either as bright lines for strong sources or dark grey ones for weak sources. Since EVLA data sets tend to be large and unwieldy, it is recommended that you separate the data into the separate target sources, applying the current calibration and flagging once and for all. The imaging task IMAGR can do this on the fly, but, especially for observations of spectral-line sources, this is excessively expensive. <source lang="text"> > DEFAULT SPLIT ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > CALCODE ’-CAL’ CR to do just target sources now. > GO CR to write out separate calibrated data sets for each target source. </source> At present, EVLA data sets have no meaningful weights associated with the visibility data. There is a task new to 31DEC10 called REWAY which computes a robust rms over spectral channels within each IF and polarization. It can simply base the weights on these on a record-by-record, baseline-by-baseline basis or it can average the rmses in time, solve for antenna-based rmses, smooth those further in time, and only then apply them to the data. For these weights to be meaningful, the bandpass calibration must be applied and any RFI or other real spectral-line signal channels must be omitted from the rms computation. For the weights to be correctly calibrated, all amplitude calibration must also be applied. For these reasons, REWAY might well be used instead of SPLIT, running it one source at a time. Thus, <source lang="text"> > DEFAULT REWAY ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > SOURCE ’target1’ , ’ ’ CR to do one target source. > GO CR to write out a calibrated, weighted data set for the first target source. </source> Then, when that finishes <source lang="text"> > SOURCE ’target2’ , ’ ’ ; GO CR to do another target source. </source> It is not clear that this algorithm is optimal, but it certainly should be better than using all weights 1.0 throughout. == Spectral-line Imaging Hints == == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] c8bb57db1ff9b9d4f23b69cc883b44fdbf2d9b56 24 23 2010-04-13T14:13:15Z Jmcmulli 2 /* Spectral-line Imaging Hints */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 CR to apply the delay calibration — very important. > SOLINT 0 CR to compute a bandpass solution for each scan on the bandpass calibrator. > CALSOUR ’bandpass cal’ CR to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. Remember these values — you will use them again. > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR to normalize the results only after the solution is found using the channels selected by ICHANSEL. > GO CR to make a bandpass (BP) table. </source> Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: <source lang="text"> > DEFAULT CALIB ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IN2DI n2; GET2N m2 CR to select the model image on disk n2 and catalog number m2. > DOCAL1;DOBAND3 CR to apply the delay and bandpass calibration — very important. > SOLINT 0 ; NMAPS 1 CR to compute a solution for each calibration scan and use the source model. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. </source> Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: <source lang="text"> > CLR2NAME ; NMAPS 0 CR to do no models. > CALSOUR ’other cal1’, ’other cal2’ CR to select the secondary calibrators by whatever names appear in your LISTR output. > GO CR to find the remaining complex gains. </source> Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: <source lang="text"> > TASK ’GETJY’ ; INP CR to set the task name without changing other adverbs. > SOURCE CALSOUR CR to select the secondary sources by the list of name you just used. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > INP CR to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. > GO CR to adjust the gains in the SN table and the fluxes in the SU (source) table. </source> Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: <source lang="text"> > DEFAULT CLCAL ; INP CR to clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > CALCODE ’*’ CR to select all calibration sources. > SNVER 2; INVERS SNVER CR to select your solution table from CALIB. Do not include the SN table from FRING a second time! > GO CR to apply SN table 2 to CL table 2, creating CL table 3. </source> Check the result using POSSM and/or VPLOT. == Target Source Data -- Edit and SPLIT == At this point, your calibration should be finished. You should now do an initial editing on the target sources, much like that done above for the calibration sources: <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all channels into one number. > CALCODE ’-CAL’ CR to do just target sources now. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> Again, remember to set it to flag all channels and IFs. You may have to select sub-windows and force the averaging to one times ∆ t to edit in detail, or perhaps the default time averaging will be beneficial. In general, the DISPLAY AMP V DIFF is a powerful way to catch bad amplitudes and phases. It will catch drop outs either as bright lines for strong sources or dark grey ones for weak sources. Since EVLA data sets tend to be large and unwieldy, it is recommended that you separate the data into the separate target sources, applying the current calibration and flagging once and for all. The imaging task IMAGR can do this on the fly, but, especially for observations of spectral-line sources, this is excessively expensive. <source lang="text"> > DEFAULT SPLIT ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > CALCODE ’-CAL’ CR to do just target sources now. > GO CR to write out separate calibrated data sets for each target source. </source> At present, EVLA data sets have no meaningful weights associated with the visibility data. There is a task new to 31DEC10 called REWAY which computes a robust rms over spectral channels within each IF and polarization. It can simply base the weights on these on a record-by-record, baseline-by-baseline basis or it can average the rmses in time, solve for antenna-based rmses, smooth those further in time, and only then apply them to the data. For these weights to be meaningful, the bandpass calibration must be applied and any RFI or other real spectral-line signal channels must be omitted from the rms computation. For the weights to be correctly calibrated, all amplitude calibration must also be applied. For these reasons, REWAY might well be used instead of SPLIT, running it one source at a time. Thus, <source lang="text"> > DEFAULT REWAY ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > SOURCE ’target1’ , ’ ’ CR to do one target source. > GO CR to write out a calibrated, weighted data set for the first target source. </source> Then, when that finishes <source lang="text"> > SOURCE ’target2’ , ’ ’ ; GO CR to do another target source. </source> It is not clear that this algorithm is optimal, but it certainly should be better than using all weights 1.0 throughout. == Spectral-line Imaging Hints == In many spectral-line observations you will now want to separate the continuum signal from the channel- dependent signals. This is discussed in some detail in §8.3. The larger number of channels from the EVLA does mean that continuum may be estimated with greater accuracy than when there were rather few channels which were both free of edge effects and spectral-line signal. The wider total bandwidth may, however, invalidate the assumption that the continuum signal at each visibility point can be represented by a polynomial of zero or first order. If there is a single dominant continuum source offset from the phase center, the assumption may be rendered valid by shifting the data with UVLSF to center the continuum source temporarily in order to subtract it. To examine this assumption and to determine which channels appear safe to use as “continuum” channels, use POSSM. <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array > BASELINE ANTEN CR and only them. > APARM 1 , 0 CR to display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels which will assist in determining channels that should be omitted. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> Note also whether the continuum appears to be a linear function of channel. If so, then use UVLSF to fit the continuum signal, writing a continuum only and a spectral-line only data set: <source lang="text"> > DEFAULT UVLSF ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for fitting > ORDER 1 CR > DOOUTPUT 1 CR > SHIFT ∆x,∆y CR > GO CR </source> the continuum. For a multi-IF data set, you will need to select the channel ranges carefully by IF. to select fitting the continuum in real and imaginary parts with a first order polynomial in channel number. UVLSF offers orders up to four, but they are not for the faint at heart and will give bad results if there are large ranges of channels left out of the fit due to line signals. to have the continuum which was fit written as a separate data set. This may be used to image the continuum. to shift the phase center to the dominant continuum source temporarily for the fitting. to run the task. Imaging the continuum output may, in addition to any scientific value of the continuum image, provide additional flagging and even self-calibration information which may be applied to the line data. If UVLSF cannot be used, flag the channels at the edges and those with spectral signals using UVFLG. Construct a continuum image with IMAGR on this flagged, spectral-line data set. Note that you might want to reduce the size of the data set with time averaging (UVAVG) and/or channel averaging (SPLIT or AVSPC) before beginning the imaging. Imaging is discussed in detail in § 5.2 through § 5.3.6 and will not be discussed here. You may find that additional editing is needed and that iterative self-calibration is of use. Be sure to copy those flags (but not the edge and spectral-signal flags) and final SN table back to the line data set. Apply them with SPLIT and then subtract the final continuum model with UVSUB. It you have had to use the spectral-index options of IMAGR, you may do the proper subtraction including these options with OOSUB rather than UVSUB. Spectral-line imaging of EVLA data will resemble that for the old VLA except for the increased number of spectral channels and the consequent increase in the data set size. Since IMAGR must read the full data set to select the data for the next channel to be imaged, it is important that the data set be small enough to fit in computer memory if at all possible. Separating the IFs into separate files will not interfere with the imaging and will help with the data set size problem: <source lang="text"> > DEFAULT UVCOP ; INP > INDI Tn; GETN Tm CR > DOWAIT 1 CR > OUTSEQ 0 ; OUTDISK INDISK CR > FORBIF=1TON;EIF=BIF;END CR > DOWAIT -1 CR </source> to reset all adverbs and choose the task. to select the calibrated target data set on disk T n and catalog number Tm. to have the task resume AIPS only after it has finished. to avoid file name issues and select the output disk. to make separate files of each of the N IFs. to turn off waiting. OSRO data sets may not need this operation and skipping the above step will simplify any continuum imaging that you may wish to do. Doing this UVCOP step on large RSRO data sets will be worth any extra trouble it may cause. Note that you could perform the separation into separate IFs before UVLSF which will speed up POSSM and UVLSF. However, the continuum output would then have to be assembled using VBGLU, which is why the steps above were shown in the present order. Spectral-line imaging is discussed in § 8.4 as well as throughout Chapter 5. With large numbers of spectral channels, you may wish to have IMAGR find appropriate Clean boxes for you. Set IM2PARM(1) through IM2PARM(6) cautiously. IM2PARM(7) controls whether the boxes of channel n are passed on to channel n + 1. The default does not pass the boxes along when autoboxing which is probably the correct decision. The end result of the imaging will be one image “cube” for each IF since each IF has to be imaged separately even with a multi-ID input data set. (If you set BIF = 1; EIF = 0 and try to image channel 103, you will actually image the average of channel 103 from each of the IFs.) To put the individual cubes together into one large cube, use MCUBE (§ 8.5.1). == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] b3028684da6622ad58a2877dfec308e117efc05f 25 24 2010-04-13T14:28:45Z Jmcmulli 2 /* Spectral-line Imaging Hints */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 CR to apply the delay calibration — very important. > SOLINT 0 CR to compute a bandpass solution for each scan on the bandpass calibrator. > CALSOUR ’bandpass cal’ CR to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. Remember these values — you will use them again. > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR to normalize the results only after the solution is found using the channels selected by ICHANSEL. > GO CR to make a bandpass (BP) table. </source> Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: <source lang="text"> > DEFAULT CALIB ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IN2DI n2; GET2N m2 CR to select the model image on disk n2 and catalog number m2. > DOCAL1;DOBAND3 CR to apply the delay and bandpass calibration — very important. > SOLINT 0 ; NMAPS 1 CR to compute a solution for each calibration scan and use the source model. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. </source> Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: <source lang="text"> > CLR2NAME ; NMAPS 0 CR to do no models. > CALSOUR ’other cal1’, ’other cal2’ CR to select the secondary calibrators by whatever names appear in your LISTR output. > GO CR to find the remaining complex gains. </source> Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: <source lang="text"> > TASK ’GETJY’ ; INP CR to set the task name without changing other adverbs. > SOURCE CALSOUR CR to select the secondary sources by the list of name you just used. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > INP CR to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. > GO CR to adjust the gains in the SN table and the fluxes in the SU (source) table. </source> Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: <source lang="text"> > DEFAULT CLCAL ; INP CR to clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > CALCODE ’*’ CR to select all calibration sources. > SNVER 2; INVERS SNVER CR to select your solution table from CALIB. Do not include the SN table from FRING a second time! > GO CR to apply SN table 2 to CL table 2, creating CL table 3. </source> Check the result using POSSM and/or VPLOT. == Target Source Data -- Edit and SPLIT == At this point, your calibration should be finished. You should now do an initial editing on the target sources, much like that done above for the calibration sources: <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all channels into one number. > CALCODE ’-CAL’ CR to do just target sources now. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> Again, remember to set it to flag all channels and IFs. You may have to select sub-windows and force the averaging to one times ∆ t to edit in detail, or perhaps the default time averaging will be beneficial. In general, the DISPLAY AMP V DIFF is a powerful way to catch bad amplitudes and phases. It will catch drop outs either as bright lines for strong sources or dark grey ones for weak sources. Since EVLA data sets tend to be large and unwieldy, it is recommended that you separate the data into the separate target sources, applying the current calibration and flagging once and for all. The imaging task IMAGR can do this on the fly, but, especially for observations of spectral-line sources, this is excessively expensive. <source lang="text"> > DEFAULT SPLIT ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > CALCODE ’-CAL’ CR to do just target sources now. > GO CR to write out separate calibrated data sets for each target source. </source> At present, EVLA data sets have no meaningful weights associated with the visibility data. There is a task new to 31DEC10 called REWAY which computes a robust rms over spectral channels within each IF and polarization. It can simply base the weights on these on a record-by-record, baseline-by-baseline basis or it can average the rmses in time, solve for antenna-based rmses, smooth those further in time, and only then apply them to the data. For these weights to be meaningful, the bandpass calibration must be applied and any RFI or other real spectral-line signal channels must be omitted from the rms computation. For the weights to be correctly calibrated, all amplitude calibration must also be applied. For these reasons, REWAY might well be used instead of SPLIT, running it one source at a time. Thus, <source lang="text"> > DEFAULT REWAY ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > SOURCE ’target1’ , ’ ’ CR to do one target source. > GO CR to write out a calibrated, weighted data set for the first target source. </source> Then, when that finishes <source lang="text"> > SOURCE ’target2’ , ’ ’ ; GO CR to do another target source. </source> It is not clear that this algorithm is optimal, but it certainly should be better than using all weights 1.0 throughout. == Spectral-line Imaging Hints == In many spectral-line observations you will now want to separate the continuum signal from the channel- dependent signals. This is discussed in some detail in §8.3. The larger number of channels from the EVLA does mean that continuum may be estimated with greater accuracy than when there were rather few channels which were both free of edge effects and spectral-line signal. The wider total bandwidth may, however, invalidate the assumption that the continuum signal at each visibility point can be represented by a polynomial of zero or first order. If there is a single dominant continuum source offset from the phase center, the assumption may be rendered valid by shifting the data with UVLSF to center the continuum source temporarily in order to subtract it. To examine this assumption and to determine which channels appear safe to use as “continuum” channels, use POSSM. <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array > BASELINE ANTEN CR and only them. > APARM 1 , 0 CR to display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels which will assist in determining channels that should be omitted. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> Note also whether the continuum appears to be a linear function of channel. If so, then use UVLSF to fit the continuum signal, writing a continuum only and a spectral-line only data set: <source lang="text"> > DEFAULT UVLSF ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for fitting the continuum. For a multi-IF data set, you will need to select the channel ranges carefully by IF. > ORDER 1 CR to select fitting the continuum in real and imaginary parts with a first order polynomial in channel number. UVLSF offers orders up to four, but they are not for the faint at heart and will give bad results if there are large ranges of channels left out of the fit due to line signals. > DOOUTPUT 1 CR to have the continuum which was fit written as a separate data set. This may be used to image the continuum. > SHIFT ∆x,∆y CR to shift the phase center to the dominant continuum source temporarily for the fitting. > GO CR to run the task. </source> Imaging the continuum output may, in addition to any scientific value of the continuum image, provide additional flagging and even self-calibration information which may be applied to the line data. If UVLSF cannot be used, flag the channels at the edges and those with spectral signals using UVFLG. Construct a continuum image with IMAGR on this flagged, spectral-line data set. Note that you might want to reduce the size of the data set with time averaging (UVAVG) and/or channel averaging (SPLIT or AVSPC) before beginning the imaging. Imaging is discussed in detail in § 5.2 through § 5.3.6 and will not be discussed here. You may find that additional editing is needed and that iterative self-calibration is of use. Be sure to copy those flags (but not the edge and spectral-signal flags) and final SN table back to the line data set. Apply them with SPLIT and then subtract the final continuum model with UVSUB. It you have had to use the spectral-index options of IMAGR, you may do the proper subtraction including these options with OOSUB rather than UVSUB. Spectral-line imaging of EVLA data will resemble that for the old VLA except for the increased number of spectral channels and the consequent increase in the data set size. Since IMAGR must read the full data set to select the data for the next channel to be imaged, it is important that the data set be small enough to fit in computer memory if at all possible. Separating the IFs into separate files will not interfere with the imaging and will help with the data set size problem: <source lang="text"> > DEFAULT UVCOP ; INP to reset all adverbs and choose the task. > INDI Tn; GETN Tm CR to select the calibrated target data set on disk T n and catalog number Tm. > DOWAIT 1 CR to have the task resume AIPS only after it has finished. > OUTSEQ 0 ; OUTDISK INDISK CR to avoid file name issues and select the output disk. > FORBIF=1TON;EIF=BIF;END CR to make separate files of each of the N IFs. > DOWAIT -1 CR to turn off waiting. </source> OSRO data sets may not need this operation and skipping the above step will simplify any continuum imaging that you may wish to do. Doing this UVCOP step on large RSRO data sets will be worth any extra trouble it may cause. Note that you could perform the separation into separate IFs before UVLSF which will speed up POSSM and UVLSF. However, the continuum output would then have to be assembled using VBGLU, which is why the steps above were shown in the present order. Spectral-line imaging is discussed in § 8.4 as well as throughout Chapter 5. With large numbers of spectral channels, you may wish to have IMAGR find appropriate Clean boxes for you. Set IM2PARM(1) through IM2PARM(6) cautiously. IM2PARM(7) controls whether the boxes of channel n are passed on to channel n + 1. The default does not pass the boxes along when autoboxing which is probably the correct decision. The end result of the imaging will be one image “cube” for each IF since each IF has to be imaged separately even with a multi-ID input data set. (If you set BIF = 1; EIF = 0 and try to image channel 103, you will actually image the average of channel 103 from each of the IFs.) To put the individual cubes together into one large cube, use MCUBE (§ 8.5.1). == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] 0c71034ec026fb8ca14b50167a84d9afa611a6dd 26 25 2010-04-13T14:34:02Z Jmcmulli 2 /* Spectral-line Imaging Hints */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 CR to apply the delay calibration — very important. > SOLINT 0 CR to compute a bandpass solution for each scan on the bandpass calibrator. > CALSOUR ’bandpass cal’ CR to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. Remember these values — you will use them again. > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR to normalize the results only after the solution is found using the channels selected by ICHANSEL. > GO CR to make a bandpass (BP) table. </source> Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: <source lang="text"> > DEFAULT CALIB ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IN2DI n2; GET2N m2 CR to select the model image on disk n2 and catalog number m2. > DOCAL1;DOBAND3 CR to apply the delay and bandpass calibration — very important. > SOLINT 0 ; NMAPS 1 CR to compute a solution for each calibration scan and use the source model. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. </source> Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: <source lang="text"> > CLR2NAME ; NMAPS 0 CR to do no models. > CALSOUR ’other cal1’, ’other cal2’ CR to select the secondary calibrators by whatever names appear in your LISTR output. > GO CR to find the remaining complex gains. </source> Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: <source lang="text"> > TASK ’GETJY’ ; INP CR to set the task name without changing other adverbs. > SOURCE CALSOUR CR to select the secondary sources by the list of name you just used. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > INP CR to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. > GO CR to adjust the gains in the SN table and the fluxes in the SU (source) table. </source> Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: <source lang="text"> > DEFAULT CLCAL ; INP CR to clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > CALCODE ’*’ CR to select all calibration sources. > SNVER 2; INVERS SNVER CR to select your solution table from CALIB. Do not include the SN table from FRING a second time! > GO CR to apply SN table 2 to CL table 2, creating CL table 3. </source> Check the result using POSSM and/or VPLOT. == Target Source Data -- Edit and SPLIT == At this point, your calibration should be finished. You should now do an initial editing on the target sources, much like that done above for the calibration sources: <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all channels into one number. > CALCODE ’-CAL’ CR to do just target sources now. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> Again, remember to set it to flag all channels and IFs. You may have to select sub-windows and force the averaging to one times ∆ t to edit in detail, or perhaps the default time averaging will be beneficial. In general, the DISPLAY AMP V DIFF is a powerful way to catch bad amplitudes and phases. It will catch drop outs either as bright lines for strong sources or dark grey ones for weak sources. Since EVLA data sets tend to be large and unwieldy, it is recommended that you separate the data into the separate target sources, applying the current calibration and flagging once and for all. The imaging task IMAGR can do this on the fly, but, especially for observations of spectral-line sources, this is excessively expensive. <source lang="text"> > DEFAULT SPLIT ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > CALCODE ’-CAL’ CR to do just target sources now. > GO CR to write out separate calibrated data sets for each target source. </source> At present, EVLA data sets have no meaningful weights associated with the visibility data. There is a task new to 31DEC10 called REWAY which computes a robust rms over spectral channels within each IF and polarization. It can simply base the weights on these on a record-by-record, baseline-by-baseline basis or it can average the rmses in time, solve for antenna-based rmses, smooth those further in time, and only then apply them to the data. For these weights to be meaningful, the bandpass calibration must be applied and any RFI or other real spectral-line signal channels must be omitted from the rms computation. For the weights to be correctly calibrated, all amplitude calibration must also be applied. For these reasons, REWAY might well be used instead of SPLIT, running it one source at a time. Thus, <source lang="text"> > DEFAULT REWAY ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > SOURCE ’target1’ , ’ ’ CR to do one target source. > GO CR to write out a calibrated, weighted data set for the first target source. </source> Then, when that finishes <source lang="text"> > SOURCE ’target2’ , ’ ’ ; GO CR to do another target source. </source> It is not clear that this algorithm is optimal, but it certainly should be better than using all weights 1.0 throughout. == Spectral-line Imaging Hints == In many spectral-line observations you will now want to separate the continuum signal from the channel- dependent signals. This is discussed in some detail in §8.3. The larger number of channels from the EVLA does mean that continuum may be estimated with greater accuracy than when there were rather few channels which were both free of edge effects and spectral-line signal. The wider total bandwidth may, however, invalidate the assumption that the continuum signal at each visibility point can be represented by a polynomial of zero or first order. If there is a single dominant continuum source offset from the phase center, the assumption may be rendered valid by shifting the data with UVLSF to center the continuum source temporarily in order to subtract it. To examine this assumption and to determine which channels appear safe to use as “continuum” channels, use POSSM. <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array > BASELINE ANTEN CR and only them. > APARM 1 , 0 CR to display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels which will assist in determining channels that should be omitted. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> Note also whether the continuum appears to be a linear function of channel. If so, then use UVLSF to fit the continuum signal, writing a continuum only and a spectral-line only data set: <source lang="text"> > DEFAULT UVLSF ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for fitting the continuum. For a multi-IF data set, you will need to select the channel ranges carefully by IF. > ORDER 1 CR to select fitting the continuum in real and imaginary parts with a first order polynomial in channel number. UVLSF offers orders up to four, but they are not for the faint at heart and will give bad results if there are large ranges of channels left out of the fit due to line signals. > DOOUTPUT 1 CR to have the continuum which was fit written as a separate data set. This may be used to image the continuum. > SHIFT ∆x,∆y CR to shift the phase center to the dominant continuum source temporarily for the fitting. > GO CR to run the task. </source> Imaging the continuum output may, in addition to any scientific value of the continuum image, provide additional flagging and even self-calibration information which may be applied to the line data. If UVLSF cannot be used, flag the channels at the edges and those with spectral signals using UVFLG. Construct a continuum image with IMAGR on this flagged, spectral-line data set. Note that you might want to reduce the size of the data set with time averaging (UVAVG) and/or channel averaging (SPLIT or AVSPC) before beginning the imaging. Imaging is discussed in detail in § 5.2 through § 5.3.6 and will not be discussed here. You may find that additional editing is needed and that iterative self-calibration is of use. Be sure to copy those flags (but not the edge and spectral-signal flags) and final SN table back to the line data set. Apply them with SPLIT and then subtract the final continuum model with UVSUB. It you have had to use the spectral-index options of IMAGR, you may do the proper subtraction including these options with OOSUB rather than UVSUB. Spectral-line imaging of EVLA data will resemble that for the old VLA except for the increased number of spectral channels and the consequent increase in the data set size. Since IMAGR must read the full data set to select the data for the next channel to be imaged, it is important that the data set be small enough to fit in computer memory if at all possible. Separating the IFs into separate files will not interfere with the imaging and will help with the data set size problem: <source lang="text"> > DEFAULT UVCOP ; INP to reset all adverbs and choose the task. > INDI Tn; GETN Tm CR to select the calibrated target data set on disk T n and catalog number Tm. > DOWAIT 1 CR to have the task resume AIPS only after it has finished. > OUTSEQ 0 ; OUTDISK INDISK CR to avoid file name issues and select the output disk. > FORBIF=1TON;EIF=BIF;END CR to make separate files of each of the N IFs. > DOWAIT -1 CR to turn off waiting. </source> OSRO data sets may not need this operation and skipping the above step will simplify any continuum imaging that you may wish to do. Doing this UVCOP step on large RSRO data sets will be worth any extra trouble it may cause. Note that you could perform the separation into separate IFs before UVLSF which will speed up POSSM and UVLSF. However, the continuum output would then have to be assembled using VBGLU, which is why the steps above were shown in the present order. Spectral-line imaging is discussed in § 8.4 as well as throughout Chapter 5. With large numbers of spectral channels, you may wish to have IMAGR find appropriate Clean boxes for you. Set IM2PARM(1) through IM2PARM(6) cautiously. IM2PARM(7) controls whether the boxes of channel n are passed on to channel n + 1. The default does not pass the boxes along when autoboxing which is probably the correct decision. The end result of the imaging will be one image “cube” for each IF since each IF has to be imaged separately even with a multi-ID input data set. (If you set BIF = 1; EIF = 0 and try to image channel 103, you will actually image the average of channel 103 from each of the IFs.) To put the individual cubes together into one large cube, use MCUBE (§ 8.5.1). == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] 0aedbdce86081b721a0256bc72e9fa81b874fd12 27 26 2010-04-13T14:37:28Z Jmcmulli 2 /* Concluding Remarks, References, Pictures */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 CR to apply the delay calibration — very important. > SOLINT 0 CR to compute a bandpass solution for each scan on the bandpass calibrator. > CALSOUR ’bandpass cal’ CR to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. Remember these values — you will use them again. > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR to normalize the results only after the solution is found using the channels selected by ICHANSEL. > GO CR to make a bandpass (BP) table. </source> Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: <source lang="text"> > DEFAULT CALIB ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IN2DI n2; GET2N m2 CR to select the model image on disk n2 and catalog number m2. > DOCAL1;DOBAND3 CR to apply the delay and bandpass calibration — very important. > SOLINT 0 ; NMAPS 1 CR to compute a solution for each calibration scan and use the source model. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. </source> Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: <source lang="text"> > CLR2NAME ; NMAPS 0 CR to do no models. > CALSOUR ’other cal1’, ’other cal2’ CR to select the secondary calibrators by whatever names appear in your LISTR output. > GO CR to find the remaining complex gains. </source> Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: <source lang="text"> > TASK ’GETJY’ ; INP CR to set the task name without changing other adverbs. > SOURCE CALSOUR CR to select the secondary sources by the list of name you just used. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > INP CR to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. > GO CR to adjust the gains in the SN table and the fluxes in the SU (source) table. </source> Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: <source lang="text"> > DEFAULT CLCAL ; INP CR to clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > CALCODE ’*’ CR to select all calibration sources. > SNVER 2; INVERS SNVER CR to select your solution table from CALIB. Do not include the SN table from FRING a second time! > GO CR to apply SN table 2 to CL table 2, creating CL table 3. </source> Check the result using POSSM and/or VPLOT. == Target Source Data -- Edit and SPLIT == At this point, your calibration should be finished. You should now do an initial editing on the target sources, much like that done above for the calibration sources: <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all channels into one number. > CALCODE ’-CAL’ CR to do just target sources now. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> Again, remember to set it to flag all channels and IFs. You may have to select sub-windows and force the averaging to one times ∆ t to edit in detail, or perhaps the default time averaging will be beneficial. In general, the DISPLAY AMP V DIFF is a powerful way to catch bad amplitudes and phases. It will catch drop outs either as bright lines for strong sources or dark grey ones for weak sources. Since EVLA data sets tend to be large and unwieldy, it is recommended that you separate the data into the separate target sources, applying the current calibration and flagging once and for all. The imaging task IMAGR can do this on the fly, but, especially for observations of spectral-line sources, this is excessively expensive. <source lang="text"> > DEFAULT SPLIT ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > CALCODE ’-CAL’ CR to do just target sources now. > GO CR to write out separate calibrated data sets for each target source. </source> At present, EVLA data sets have no meaningful weights associated with the visibility data. There is a task new to 31DEC10 called REWAY which computes a robust rms over spectral channels within each IF and polarization. It can simply base the weights on these on a record-by-record, baseline-by-baseline basis or it can average the rmses in time, solve for antenna-based rmses, smooth those further in time, and only then apply them to the data. For these weights to be meaningful, the bandpass calibration must be applied and any RFI or other real spectral-line signal channels must be omitted from the rms computation. For the weights to be correctly calibrated, all amplitude calibration must also be applied. For these reasons, REWAY might well be used instead of SPLIT, running it one source at a time. Thus, <source lang="text"> > DEFAULT REWAY ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > SOURCE ’target1’ , ’ ’ CR to do one target source. > GO CR to write out a calibrated, weighted data set for the first target source. </source> Then, when that finishes <source lang="text"> > SOURCE ’target2’ , ’ ’ ; GO CR to do another target source. </source> It is not clear that this algorithm is optimal, but it certainly should be better than using all weights 1.0 throughout. == Spectral-line Imaging Hints == In many spectral-line observations you will now want to separate the continuum signal from the channel- dependent signals. This is discussed in some detail in §8.3. The larger number of channels from the EVLA does mean that continuum may be estimated with greater accuracy than when there were rather few channels which were both free of edge effects and spectral-line signal. The wider total bandwidth may, however, invalidate the assumption that the continuum signal at each visibility point can be represented by a polynomial of zero or first order. If there is a single dominant continuum source offset from the phase center, the assumption may be rendered valid by shifting the data with UVLSF to center the continuum source temporarily in order to subtract it. To examine this assumption and to determine which channels appear safe to use as “continuum” channels, use POSSM. <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array > BASELINE ANTEN CR and only them. > APARM 1 , 0 CR to display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels which will assist in determining channels that should be omitted. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> Note also whether the continuum appears to be a linear function of channel. If so, then use UVLSF to fit the continuum signal, writing a continuum only and a spectral-line only data set: <source lang="text"> > DEFAULT UVLSF ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for fitting the continuum. For a multi-IF data set, you will need to select the channel ranges carefully by IF. > ORDER 1 CR to select fitting the continuum in real and imaginary parts with a first order polynomial in channel number. UVLSF offers orders up to four, but they are not for the faint at heart and will give bad results if there are large ranges of channels left out of the fit due to line signals. > DOOUTPUT 1 CR to have the continuum which was fit written as a separate data set. This may be used to image the continuum. > SHIFT ∆x,∆y CR to shift the phase center to the dominant continuum source temporarily for the fitting. > GO CR to run the task. </source> Imaging the continuum output may, in addition to any scientific value of the continuum image, provide additional flagging and even self-calibration information which may be applied to the line data. If UVLSF cannot be used, flag the channels at the edges and those with spectral signals using UVFLG. Construct a continuum image with IMAGR on this flagged, spectral-line data set. Note that you might want to reduce the size of the data set with time averaging (UVAVG) and/or channel averaging (SPLIT or AVSPC) before beginning the imaging. Imaging is discussed in detail in § 5.2 through § 5.3.6 and will not be discussed here. You may find that additional editing is needed and that iterative self-calibration is of use. Be sure to copy those flags (but not the edge and spectral-signal flags) and final SN table back to the line data set. Apply them with SPLIT and then subtract the final continuum model with UVSUB. It you have had to use the spectral-index options of IMAGR, you may do the proper subtraction including these options with OOSUB rather than UVSUB. Spectral-line imaging of EVLA data will resemble that for the old VLA except for the increased number of spectral channels and the consequent increase in the data set size. Since IMAGR must read the full data set to select the data for the next channel to be imaged, it is important that the data set be small enough to fit in computer memory if at all possible. Separating the IFs into separate files will not interfere with the imaging and will help with the data set size problem: <source lang="text"> > DEFAULT UVCOP ; INP to reset all adverbs and choose the task. > INDI Tn; GETN Tm CR to select the calibrated target data set on disk T n and catalog number Tm. > DOWAIT 1 CR to have the task resume AIPS only after it has finished. > OUTSEQ 0 ; OUTDISK INDISK CR to avoid file name issues and select the output disk. > FORBIF=1TON;EIF=BIF;END CR to make separate files of each of the N IFs. > DOWAIT -1 CR to turn off waiting. </source> OSRO data sets may not need this operation and skipping the above step will simplify any continuum imaging that you may wish to do. Doing this UVCOP step on large RSRO data sets will be worth any extra trouble it may cause. Note that you could perform the separation into separate IFs before UVLSF which will speed up POSSM and UVLSF. However, the continuum output would then have to be assembled using VBGLU, which is why the steps above were shown in the present order. Spectral-line imaging is discussed in § 8.4 as well as throughout Chapter 5. With large numbers of spectral channels, you may wish to have IMAGR find appropriate Clean boxes for you. Set IM2PARM(1) through IM2PARM(6) cautiously. IM2PARM(7) controls whether the boxes of channel n are passed on to channel n + 1. The default does not pass the boxes along when autoboxing which is probably the correct decision. The end result of the imaging will be one image “cube” for each IF since each IF has to be imaged separately even with a multi-ID input data set. (If you set BIF = 1; EIF = 0 and try to image channel 103, you will actually image the average of channel 103 from each of the IFs.) To put the individual cubes together into one large cube, use MCUBE (§ 8.5.1). == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] Figure E.1: The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten. 5378722cd91e23b50ec55243c7020b192c916234 28 27 2010-04-13T14:40:52Z Jmcmulli 2 /* Overview */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. '''Note:''' This document is written as an appendix to the AIPS Cookbook; section numbers refer to that document: http://www.aips.nrao.edu/cook.html == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 CR to apply the delay calibration — very important. > SOLINT 0 CR to compute a bandpass solution for each scan on the bandpass calibrator. > CALSOUR ’bandpass cal’ CR to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. Remember these values — you will use them again. > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR to normalize the results only after the solution is found using the channels selected by ICHANSEL. > GO CR to make a bandpass (BP) table. </source> Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: <source lang="text"> > DEFAULT CALIB ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IN2DI n2; GET2N m2 CR to select the model image on disk n2 and catalog number m2. > DOCAL1;DOBAND3 CR to apply the delay and bandpass calibration — very important. > SOLINT 0 ; NMAPS 1 CR to compute a solution for each calibration scan and use the source model. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. </source> Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: <source lang="text"> > CLR2NAME ; NMAPS 0 CR to do no models. > CALSOUR ’other cal1’, ’other cal2’ CR to select the secondary calibrators by whatever names appear in your LISTR output. > GO CR to find the remaining complex gains. </source> Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: <source lang="text"> > TASK ’GETJY’ ; INP CR to set the task name without changing other adverbs. > SOURCE CALSOUR CR to select the secondary sources by the list of name you just used. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > INP CR to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. > GO CR to adjust the gains in the SN table and the fluxes in the SU (source) table. </source> Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: <source lang="text"> > DEFAULT CLCAL ; INP CR to clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > CALCODE ’*’ CR to select all calibration sources. > SNVER 2; INVERS SNVER CR to select your solution table from CALIB. Do not include the SN table from FRING a second time! > GO CR to apply SN table 2 to CL table 2, creating CL table 3. </source> Check the result using POSSM and/or VPLOT. == Target Source Data -- Edit and SPLIT == At this point, your calibration should be finished. You should now do an initial editing on the target sources, much like that done above for the calibration sources: <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all channels into one number. > CALCODE ’-CAL’ CR to do just target sources now. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> Again, remember to set it to flag all channels and IFs. You may have to select sub-windows and force the averaging to one times ∆ t to edit in detail, or perhaps the default time averaging will be beneficial. In general, the DISPLAY AMP V DIFF is a powerful way to catch bad amplitudes and phases. It will catch drop outs either as bright lines for strong sources or dark grey ones for weak sources. Since EVLA data sets tend to be large and unwieldy, it is recommended that you separate the data into the separate target sources, applying the current calibration and flagging once and for all. The imaging task IMAGR can do this on the fly, but, especially for observations of spectral-line sources, this is excessively expensive. <source lang="text"> > DEFAULT SPLIT ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > CALCODE ’-CAL’ CR to do just target sources now. > GO CR to write out separate calibrated data sets for each target source. </source> At present, EVLA data sets have no meaningful weights associated with the visibility data. There is a task new to 31DEC10 called REWAY which computes a robust rms over spectral channels within each IF and polarization. It can simply base the weights on these on a record-by-record, baseline-by-baseline basis or it can average the rmses in time, solve for antenna-based rmses, smooth those further in time, and only then apply them to the data. For these weights to be meaningful, the bandpass calibration must be applied and any RFI or other real spectral-line signal channels must be omitted from the rms computation. For the weights to be correctly calibrated, all amplitude calibration must also be applied. For these reasons, REWAY might well be used instead of SPLIT, running it one source at a time. Thus, <source lang="text"> > DEFAULT REWAY ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > SOURCE ’target1’ , ’ ’ CR to do one target source. > GO CR to write out a calibrated, weighted data set for the first target source. </source> Then, when that finishes <source lang="text"> > SOURCE ’target2’ , ’ ’ ; GO CR to do another target source. </source> It is not clear that this algorithm is optimal, but it certainly should be better than using all weights 1.0 throughout. == Spectral-line Imaging Hints == In many spectral-line observations you will now want to separate the continuum signal from the channel- dependent signals. This is discussed in some detail in §8.3. The larger number of channels from the EVLA does mean that continuum may be estimated with greater accuracy than when there were rather few channels which were both free of edge effects and spectral-line signal. The wider total bandwidth may, however, invalidate the assumption that the continuum signal at each visibility point can be represented by a polynomial of zero or first order. If there is a single dominant continuum source offset from the phase center, the assumption may be rendered valid by shifting the data with UVLSF to center the continuum source temporarily in order to subtract it. To examine this assumption and to determine which channels appear safe to use as “continuum” channels, use POSSM. <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array > BASELINE ANTEN CR and only them. > APARM 1 , 0 CR to display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels which will assist in determining channels that should be omitted. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> Note also whether the continuum appears to be a linear function of channel. If so, then use UVLSF to fit the continuum signal, writing a continuum only and a spectral-line only data set: <source lang="text"> > DEFAULT UVLSF ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for fitting the continuum. For a multi-IF data set, you will need to select the channel ranges carefully by IF. > ORDER 1 CR to select fitting the continuum in real and imaginary parts with a first order polynomial in channel number. UVLSF offers orders up to four, but they are not for the faint at heart and will give bad results if there are large ranges of channels left out of the fit due to line signals. > DOOUTPUT 1 CR to have the continuum which was fit written as a separate data set. This may be used to image the continuum. > SHIFT ∆x,∆y CR to shift the phase center to the dominant continuum source temporarily for the fitting. > GO CR to run the task. </source> Imaging the continuum output may, in addition to any scientific value of the continuum image, provide additional flagging and even self-calibration information which may be applied to the line data. If UVLSF cannot be used, flag the channels at the edges and those with spectral signals using UVFLG. Construct a continuum image with IMAGR on this flagged, spectral-line data set. Note that you might want to reduce the size of the data set with time averaging (UVAVG) and/or channel averaging (SPLIT or AVSPC) before beginning the imaging. Imaging is discussed in detail in § 5.2 through § 5.3.6 and will not be discussed here. You may find that additional editing is needed and that iterative self-calibration is of use. Be sure to copy those flags (but not the edge and spectral-signal flags) and final SN table back to the line data set. Apply them with SPLIT and then subtract the final continuum model with UVSUB. It you have had to use the spectral-index options of IMAGR, you may do the proper subtraction including these options with OOSUB rather than UVSUB. Spectral-line imaging of EVLA data will resemble that for the old VLA except for the increased number of spectral channels and the consequent increase in the data set size. Since IMAGR must read the full data set to select the data for the next channel to be imaged, it is important that the data set be small enough to fit in computer memory if at all possible. Separating the IFs into separate files will not interfere with the imaging and will help with the data set size problem: <source lang="text"> > DEFAULT UVCOP ; INP to reset all adverbs and choose the task. > INDI Tn; GETN Tm CR to select the calibrated target data set on disk T n and catalog number Tm. > DOWAIT 1 CR to have the task resume AIPS only after it has finished. > OUTSEQ 0 ; OUTDISK INDISK CR to avoid file name issues and select the output disk. > FORBIF=1TON;EIF=BIF;END CR to make separate files of each of the N IFs. > DOWAIT -1 CR to turn off waiting. </source> OSRO data sets may not need this operation and skipping the above step will simplify any continuum imaging that you may wish to do. Doing this UVCOP step on large RSRO data sets will be worth any extra trouble it may cause. Note that you could perform the separation into separate IFs before UVLSF which will speed up POSSM and UVLSF. However, the continuum output would then have to be assembled using VBGLU, which is why the steps above were shown in the present order. Spectral-line imaging is discussed in § 8.4 as well as throughout Chapter 5. With large numbers of spectral channels, you may wish to have IMAGR find appropriate Clean boxes for you. Set IM2PARM(1) through IM2PARM(6) cautiously. IM2PARM(7) controls whether the boxes of channel n are passed on to channel n + 1. The default does not pass the boxes along when autoboxing which is probably the correct decision. The end result of the imaging will be one image “cube” for each IF since each IF has to be imaged separately even with a multi-ID input data set. (If you set BIF = 1; EIF = 0 and try to image channel 103, you will actually image the average of channel 103 from each of the IFs.) To put the individual cubes together into one large cube, use MCUBE (§ 8.5.1). == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] Figure E.1: The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten. 76b612d2b9fcc6c6ac1a9c2927af9acad9045142 29 28 2010-04-13T14:43:04Z Jmcmulli 2 /* Getting Your Data into AIPS */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. '''Note:''' This document is written as an appendix to the AIPS Cookbook; section numbers refer to that document: http://www.aips.nrao.edu/cook.html == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 CR to apply the delay calibration — very important. > SOLINT 0 CR to compute a bandpass solution for each scan on the bandpass calibrator. > CALSOUR ’bandpass cal’ CR to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. Remember these values — you will use them again. > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR to normalize the results only after the solution is found using the channels selected by ICHANSEL. > GO CR to make a bandpass (BP) table. </source> Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: <source lang="text"> > DEFAULT CALIB ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IN2DI n2; GET2N m2 CR to select the model image on disk n2 and catalog number m2. > DOCAL1;DOBAND3 CR to apply the delay and bandpass calibration — very important. > SOLINT 0 ; NMAPS 1 CR to compute a solution for each calibration scan and use the source model. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. </source> Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: <source lang="text"> > CLR2NAME ; NMAPS 0 CR to do no models. > CALSOUR ’other cal1’, ’other cal2’ CR to select the secondary calibrators by whatever names appear in your LISTR output. > GO CR to find the remaining complex gains. </source> Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: <source lang="text"> > TASK ’GETJY’ ; INP CR to set the task name without changing other adverbs. > SOURCE CALSOUR CR to select the secondary sources by the list of name you just used. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > INP CR to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. > GO CR to adjust the gains in the SN table and the fluxes in the SU (source) table. </source> Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: <source lang="text"> > DEFAULT CLCAL ; INP CR to clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > CALCODE ’*’ CR to select all calibration sources. > SNVER 2; INVERS SNVER CR to select your solution table from CALIB. Do not include the SN table from FRING a second time! > GO CR to apply SN table 2 to CL table 2, creating CL table 3. </source> Check the result using POSSM and/or VPLOT. == Target Source Data -- Edit and SPLIT == At this point, your calibration should be finished. You should now do an initial editing on the target sources, much like that done above for the calibration sources: <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all channels into one number. > CALCODE ’-CAL’ CR to do just target sources now. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> Again, remember to set it to flag all channels and IFs. You may have to select sub-windows and force the averaging to one times ∆ t to edit in detail, or perhaps the default time averaging will be beneficial. In general, the DISPLAY AMP V DIFF is a powerful way to catch bad amplitudes and phases. It will catch drop outs either as bright lines for strong sources or dark grey ones for weak sources. Since EVLA data sets tend to be large and unwieldy, it is recommended that you separate the data into the separate target sources, applying the current calibration and flagging once and for all. The imaging task IMAGR can do this on the fly, but, especially for observations of spectral-line sources, this is excessively expensive. <source lang="text"> > DEFAULT SPLIT ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > CALCODE ’-CAL’ CR to do just target sources now. > GO CR to write out separate calibrated data sets for each target source. </source> At present, EVLA data sets have no meaningful weights associated with the visibility data. There is a task new to 31DEC10 called REWAY which computes a robust rms over spectral channels within each IF and polarization. It can simply base the weights on these on a record-by-record, baseline-by-baseline basis or it can average the rmses in time, solve for antenna-based rmses, smooth those further in time, and only then apply them to the data. For these weights to be meaningful, the bandpass calibration must be applied and any RFI or other real spectral-line signal channels must be omitted from the rms computation. For the weights to be correctly calibrated, all amplitude calibration must also be applied. For these reasons, REWAY might well be used instead of SPLIT, running it one source at a time. Thus, <source lang="text"> > DEFAULT REWAY ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > SOURCE ’target1’ , ’ ’ CR to do one target source. > GO CR to write out a calibrated, weighted data set for the first target source. </source> Then, when that finishes <source lang="text"> > SOURCE ’target2’ , ’ ’ ; GO CR to do another target source. </source> It is not clear that this algorithm is optimal, but it certainly should be better than using all weights 1.0 throughout. == Spectral-line Imaging Hints == In many spectral-line observations you will now want to separate the continuum signal from the channel- dependent signals. This is discussed in some detail in §8.3. The larger number of channels from the EVLA does mean that continuum may be estimated with greater accuracy than when there were rather few channels which were both free of edge effects and spectral-line signal. The wider total bandwidth may, however, invalidate the assumption that the continuum signal at each visibility point can be represented by a polynomial of zero or first order. If there is a single dominant continuum source offset from the phase center, the assumption may be rendered valid by shifting the data with UVLSF to center the continuum source temporarily in order to subtract it. To examine this assumption and to determine which channels appear safe to use as “continuum” channels, use POSSM. <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array > BASELINE ANTEN CR and only them. > APARM 1 , 0 CR to display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels which will assist in determining channels that should be omitted. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> Note also whether the continuum appears to be a linear function of channel. If so, then use UVLSF to fit the continuum signal, writing a continuum only and a spectral-line only data set: <source lang="text"> > DEFAULT UVLSF ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for fitting the continuum. For a multi-IF data set, you will need to select the channel ranges carefully by IF. > ORDER 1 CR to select fitting the continuum in real and imaginary parts with a first order polynomial in channel number. UVLSF offers orders up to four, but they are not for the faint at heart and will give bad results if there are large ranges of channels left out of the fit due to line signals. > DOOUTPUT 1 CR to have the continuum which was fit written as a separate data set. This may be used to image the continuum. > SHIFT ∆x,∆y CR to shift the phase center to the dominant continuum source temporarily for the fitting. > GO CR to run the task. </source> Imaging the continuum output may, in addition to any scientific value of the continuum image, provide additional flagging and even self-calibration information which may be applied to the line data. If UVLSF cannot be used, flag the channels at the edges and those with spectral signals using UVFLG. Construct a continuum image with IMAGR on this flagged, spectral-line data set. Note that you might want to reduce the size of the data set with time averaging (UVAVG) and/or channel averaging (SPLIT or AVSPC) before beginning the imaging. Imaging is discussed in detail in § 5.2 through § 5.3.6 and will not be discussed here. You may find that additional editing is needed and that iterative self-calibration is of use. Be sure to copy those flags (but not the edge and spectral-signal flags) and final SN table back to the line data set. Apply them with SPLIT and then subtract the final continuum model with UVSUB. It you have had to use the spectral-index options of IMAGR, you may do the proper subtraction including these options with OOSUB rather than UVSUB. Spectral-line imaging of EVLA data will resemble that for the old VLA except for the increased number of spectral channels and the consequent increase in the data set size. Since IMAGR must read the full data set to select the data for the next channel to be imaged, it is important that the data set be small enough to fit in computer memory if at all possible. Separating the IFs into separate files will not interfere with the imaging and will help with the data set size problem: <source lang="text"> > DEFAULT UVCOP ; INP to reset all adverbs and choose the task. > INDI Tn; GETN Tm CR to select the calibrated target data set on disk T n and catalog number Tm. > DOWAIT 1 CR to have the task resume AIPS only after it has finished. > OUTSEQ 0 ; OUTDISK INDISK CR to avoid file name issues and select the output disk. > FORBIF=1TON;EIF=BIF;END CR to make separate files of each of the N IFs. > DOWAIT -1 CR to turn off waiting. </source> OSRO data sets may not need this operation and skipping the above step will simplify any continuum imaging that you may wish to do. Doing this UVCOP step on large RSRO data sets will be worth any extra trouble it may cause. Note that you could perform the separation into separate IFs before UVLSF which will speed up POSSM and UVLSF. However, the continuum output would then have to be assembled using VBGLU, which is why the steps above were shown in the present order. Spectral-line imaging is discussed in § 8.4 as well as throughout Chapter 5. With large numbers of spectral channels, you may wish to have IMAGR find appropriate Clean boxes for you. Set IM2PARM(1) through IM2PARM(6) cautiously. IM2PARM(7) controls whether the boxes of channel n are passed on to channel n + 1. The default does not pass the boxes along when autoboxing which is probably the correct decision. The end result of the imaging will be one image “cube” for each IF since each IF has to be imaged separately even with a multi-ID input data set. (If you set BIF = 1; EIF = 0 and try to image channel 103, you will actually image the average of channel 103 from each of the IFs.) To put the individual cubes together into one large cube, use MCUBE (§ 8.5.1). == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] Figure E.1: The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten. 197728c123ec6a033cd4ef83fe5e420e0acda826 30 29 2010-04-13T14:45:15Z Jmcmulli 2 /* Initial Calibration -- FRING */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. '''Note:''' This document is written as an appendix to the AIPS Cookbook; section numbers refer to that document: http://www.aips.nrao.edu/cook.html == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 CR to apply the delay calibration — very important. > SOLINT 0 CR to compute a bandpass solution for each scan on the bandpass calibrator. > CALSOUR ’bandpass cal’ CR to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. Remember these values — you will use them again. > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR to normalize the results only after the solution is found using the channels selected by ICHANSEL. > GO CR to make a bandpass (BP) table. </source> Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: <source lang="text"> > DEFAULT CALIB ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IN2DI n2; GET2N m2 CR to select the model image on disk n2 and catalog number m2. > DOCAL1;DOBAND3 CR to apply the delay and bandpass calibration — very important. > SOLINT 0 ; NMAPS 1 CR to compute a solution for each calibration scan and use the source model. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. </source> Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: <source lang="text"> > CLR2NAME ; NMAPS 0 CR to do no models. > CALSOUR ’other cal1’, ’other cal2’ CR to select the secondary calibrators by whatever names appear in your LISTR output. > GO CR to find the remaining complex gains. </source> Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: <source lang="text"> > TASK ’GETJY’ ; INP CR to set the task name without changing other adverbs. > SOURCE CALSOUR CR to select the secondary sources by the list of name you just used. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > INP CR to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. > GO CR to adjust the gains in the SN table and the fluxes in the SU (source) table. </source> Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: <source lang="text"> > DEFAULT CLCAL ; INP CR to clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > CALCODE ’*’ CR to select all calibration sources. > SNVER 2; INVERS SNVER CR to select your solution table from CALIB. Do not include the SN table from FRING a second time! > GO CR to apply SN table 2 to CL table 2, creating CL table 3. </source> Check the result using POSSM and/or VPLOT. == Target Source Data -- Edit and SPLIT == At this point, your calibration should be finished. You should now do an initial editing on the target sources, much like that done above for the calibration sources: <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all channels into one number. > CALCODE ’-CAL’ CR to do just target sources now. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> Again, remember to set it to flag all channels and IFs. You may have to select sub-windows and force the averaging to one times ∆ t to edit in detail, or perhaps the default time averaging will be beneficial. In general, the DISPLAY AMP V DIFF is a powerful way to catch bad amplitudes and phases. It will catch drop outs either as bright lines for strong sources or dark grey ones for weak sources. Since EVLA data sets tend to be large and unwieldy, it is recommended that you separate the data into the separate target sources, applying the current calibration and flagging once and for all. The imaging task IMAGR can do this on the fly, but, especially for observations of spectral-line sources, this is excessively expensive. <source lang="text"> > DEFAULT SPLIT ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > CALCODE ’-CAL’ CR to do just target sources now. > GO CR to write out separate calibrated data sets for each target source. </source> At present, EVLA data sets have no meaningful weights associated with the visibility data. There is a task new to 31DEC10 called REWAY which computes a robust rms over spectral channels within each IF and polarization. It can simply base the weights on these on a record-by-record, baseline-by-baseline basis or it can average the rmses in time, solve for antenna-based rmses, smooth those further in time, and only then apply them to the data. For these weights to be meaningful, the bandpass calibration must be applied and any RFI or other real spectral-line signal channels must be omitted from the rms computation. For the weights to be correctly calibrated, all amplitude calibration must also be applied. For these reasons, REWAY might well be used instead of SPLIT, running it one source at a time. Thus, <source lang="text"> > DEFAULT REWAY ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > SOURCE ’target1’ , ’ ’ CR to do one target source. > GO CR to write out a calibrated, weighted data set for the first target source. </source> Then, when that finishes <source lang="text"> > SOURCE ’target2’ , ’ ’ ; GO CR to do another target source. </source> It is not clear that this algorithm is optimal, but it certainly should be better than using all weights 1.0 throughout. == Spectral-line Imaging Hints == In many spectral-line observations you will now want to separate the continuum signal from the channel- dependent signals. This is discussed in some detail in §8.3. The larger number of channels from the EVLA does mean that continuum may be estimated with greater accuracy than when there were rather few channels which were both free of edge effects and spectral-line signal. The wider total bandwidth may, however, invalidate the assumption that the continuum signal at each visibility point can be represented by a polynomial of zero or first order. If there is a single dominant continuum source offset from the phase center, the assumption may be rendered valid by shifting the data with UVLSF to center the continuum source temporarily in order to subtract it. To examine this assumption and to determine which channels appear safe to use as “continuum” channels, use POSSM. <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array > BASELINE ANTEN CR and only them. > APARM 1 , 0 CR to display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels which will assist in determining channels that should be omitted. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> Note also whether the continuum appears to be a linear function of channel. If so, then use UVLSF to fit the continuum signal, writing a continuum only and a spectral-line only data set: <source lang="text"> > DEFAULT UVLSF ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for fitting the continuum. For a multi-IF data set, you will need to select the channel ranges carefully by IF. > ORDER 1 CR to select fitting the continuum in real and imaginary parts with a first order polynomial in channel number. UVLSF offers orders up to four, but they are not for the faint at heart and will give bad results if there are large ranges of channels left out of the fit due to line signals. > DOOUTPUT 1 CR to have the continuum which was fit written as a separate data set. This may be used to image the continuum. > SHIFT ∆x,∆y CR to shift the phase center to the dominant continuum source temporarily for the fitting. > GO CR to run the task. </source> Imaging the continuum output may, in addition to any scientific value of the continuum image, provide additional flagging and even self-calibration information which may be applied to the line data. If UVLSF cannot be used, flag the channels at the edges and those with spectral signals using UVFLG. Construct a continuum image with IMAGR on this flagged, spectral-line data set. Note that you might want to reduce the size of the data set with time averaging (UVAVG) and/or channel averaging (SPLIT or AVSPC) before beginning the imaging. Imaging is discussed in detail in § 5.2 through § 5.3.6 and will not be discussed here. You may find that additional editing is needed and that iterative self-calibration is of use. Be sure to copy those flags (but not the edge and spectral-signal flags) and final SN table back to the line data set. Apply them with SPLIT and then subtract the final continuum model with UVSUB. It you have had to use the spectral-index options of IMAGR, you may do the proper subtraction including these options with OOSUB rather than UVSUB. Spectral-line imaging of EVLA data will resemble that for the old VLA except for the increased number of spectral channels and the consequent increase in the data set size. Since IMAGR must read the full data set to select the data for the next channel to be imaged, it is important that the data set be small enough to fit in computer memory if at all possible. Separating the IFs into separate files will not interfere with the imaging and will help with the data set size problem: <source lang="text"> > DEFAULT UVCOP ; INP to reset all adverbs and choose the task. > INDI Tn; GETN Tm CR to select the calibrated target data set on disk T n and catalog number Tm. > DOWAIT 1 CR to have the task resume AIPS only after it has finished. > OUTSEQ 0 ; OUTDISK INDISK CR to avoid file name issues and select the output disk. > FORBIF=1TON;EIF=BIF;END CR to make separate files of each of the N IFs. > DOWAIT -1 CR to turn off waiting. </source> OSRO data sets may not need this operation and skipping the above step will simplify any continuum imaging that you may wish to do. Doing this UVCOP step on large RSRO data sets will be worth any extra trouble it may cause. Note that you could perform the separation into separate IFs before UVLSF which will speed up POSSM and UVLSF. However, the continuum output would then have to be assembled using VBGLU, which is why the steps above were shown in the present order. Spectral-line imaging is discussed in § 8.4 as well as throughout Chapter 5. With large numbers of spectral channels, you may wish to have IMAGR find appropriate Clean boxes for you. Set IM2PARM(1) through IM2PARM(6) cautiously. IM2PARM(7) controls whether the boxes of channel n are passed on to channel n + 1. The default does not pass the boxes along when autoboxing which is probably the correct decision. The end result of the imaging will be one image “cube” for each IF since each IF has to be imaged separately even with a multi-ID input data set. (If you set BIF = 1; EIF = 0 and try to image channel 103, you will actually image the average of channel 103 from each of the IFs.) To put the individual cubes together into one large cube, use MCUBE (§ 8.5.1). == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] Figure E.1: The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten. 2e53eeb168d47eede34cff08da15bf46b76c5c7f 31 30 2010-04-13T14:46:51Z Jmcmulli 2 /* Initial Calibration -- FRING */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. '''Note:''' This document is written as an appendix to the AIPS Cookbook; section numbers refer to that document: http://www.aips.nrao.edu/cook.html == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 CR to apply the delay calibration — very important. > SOLINT 0 CR to compute a bandpass solution for each scan on the bandpass calibrator. > CALSOUR ’bandpass cal’ CR to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. Remember these values — you will use them again. > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR to normalize the results only after the solution is found using the channels selected by ICHANSEL. > GO CR to make a bandpass (BP) table. </source> Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: <source lang="text"> > DEFAULT CALIB ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IN2DI n2; GET2N m2 CR to select the model image on disk n2 and catalog number m2. > DOCAL1;DOBAND3 CR to apply the delay and bandpass calibration — very important. > SOLINT 0 ; NMAPS 1 CR to compute a solution for each calibration scan and use the source model. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. </source> Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: <source lang="text"> > CLR2NAME ; NMAPS 0 CR to do no models. > CALSOUR ’other cal1’, ’other cal2’ CR to select the secondary calibrators by whatever names appear in your LISTR output. > GO CR to find the remaining complex gains. </source> Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: <source lang="text"> > TASK ’GETJY’ ; INP CR to set the task name without changing other adverbs. > SOURCE CALSOUR CR to select the secondary sources by the list of name you just used. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > INP CR to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. > GO CR to adjust the gains in the SN table and the fluxes in the SU (source) table. </source> Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: <source lang="text"> > DEFAULT CLCAL ; INP CR to clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > CALCODE ’*’ CR to select all calibration sources. > SNVER 2; INVERS SNVER CR to select your solution table from CALIB. Do not include the SN table from FRING a second time! > GO CR to apply SN table 2 to CL table 2, creating CL table 3. </source> Check the result using POSSM and/or VPLOT. == Target Source Data -- Edit and SPLIT == At this point, your calibration should be finished. You should now do an initial editing on the target sources, much like that done above for the calibration sources: <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all channels into one number. > CALCODE ’-CAL’ CR to do just target sources now. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> Again, remember to set it to flag all channels and IFs. You may have to select sub-windows and force the averaging to one times ∆ t to edit in detail, or perhaps the default time averaging will be beneficial. In general, the DISPLAY AMP V DIFF is a powerful way to catch bad amplitudes and phases. It will catch drop outs either as bright lines for strong sources or dark grey ones for weak sources. Since EVLA data sets tend to be large and unwieldy, it is recommended that you separate the data into the separate target sources, applying the current calibration and flagging once and for all. The imaging task IMAGR can do this on the fly, but, especially for observations of spectral-line sources, this is excessively expensive. <source lang="text"> > DEFAULT SPLIT ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > CALCODE ’-CAL’ CR to do just target sources now. > GO CR to write out separate calibrated data sets for each target source. </source> At present, EVLA data sets have no meaningful weights associated with the visibility data. There is a task new to 31DEC10 called REWAY which computes a robust rms over spectral channels within each IF and polarization. It can simply base the weights on these on a record-by-record, baseline-by-baseline basis or it can average the rmses in time, solve for antenna-based rmses, smooth those further in time, and only then apply them to the data. For these weights to be meaningful, the bandpass calibration must be applied and any RFI or other real spectral-line signal channels must be omitted from the rms computation. For the weights to be correctly calibrated, all amplitude calibration must also be applied. For these reasons, REWAY might well be used instead of SPLIT, running it one source at a time. Thus, <source lang="text"> > DEFAULT REWAY ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > SOURCE ’target1’ , ’ ’ CR to do one target source. > GO CR to write out a calibrated, weighted data set for the first target source. </source> Then, when that finishes <source lang="text"> > SOURCE ’target2’ , ’ ’ ; GO CR to do another target source. </source> It is not clear that this algorithm is optimal, but it certainly should be better than using all weights 1.0 throughout. == Spectral-line Imaging Hints == In many spectral-line observations you will now want to separate the continuum signal from the channel- dependent signals. This is discussed in some detail in §8.3. The larger number of channels from the EVLA does mean that continuum may be estimated with greater accuracy than when there were rather few channels which were both free of edge effects and spectral-line signal. The wider total bandwidth may, however, invalidate the assumption that the continuum signal at each visibility point can be represented by a polynomial of zero or first order. If there is a single dominant continuum source offset from the phase center, the assumption may be rendered valid by shifting the data with UVLSF to center the continuum source temporarily in order to subtract it. To examine this assumption and to determine which channels appear safe to use as “continuum” channels, use POSSM. <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array > BASELINE ANTEN CR and only them. > APARM 1 , 0 CR to display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels which will assist in determining channels that should be omitted. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> Note also whether the continuum appears to be a linear function of channel. If so, then use UVLSF to fit the continuum signal, writing a continuum only and a spectral-line only data set: <source lang="text"> > DEFAULT UVLSF ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for fitting the continuum. For a multi-IF data set, you will need to select the channel ranges carefully by IF. > ORDER 1 CR to select fitting the continuum in real and imaginary parts with a first order polynomial in channel number. UVLSF offers orders up to four, but they are not for the faint at heart and will give bad results if there are large ranges of channels left out of the fit due to line signals. > DOOUTPUT 1 CR to have the continuum which was fit written as a separate data set. This may be used to image the continuum. > SHIFT ∆x,∆y CR to shift the phase center to the dominant continuum source temporarily for the fitting. > GO CR to run the task. </source> Imaging the continuum output may, in addition to any scientific value of the continuum image, provide additional flagging and even self-calibration information which may be applied to the line data. If UVLSF cannot be used, flag the channels at the edges and those with spectral signals using UVFLG. Construct a continuum image with IMAGR on this flagged, spectral-line data set. Note that you might want to reduce the size of the data set with time averaging (UVAVG) and/or channel averaging (SPLIT or AVSPC) before beginning the imaging. Imaging is discussed in detail in § 5.2 through § 5.3.6 and will not be discussed here. You may find that additional editing is needed and that iterative self-calibration is of use. Be sure to copy those flags (but not the edge and spectral-signal flags) and final SN table back to the line data set. Apply them with SPLIT and then subtract the final continuum model with UVSUB. It you have had to use the spectral-index options of IMAGR, you may do the proper subtraction including these options with OOSUB rather than UVSUB. Spectral-line imaging of EVLA data will resemble that for the old VLA except for the increased number of spectral channels and the consequent increase in the data set size. Since IMAGR must read the full data set to select the data for the next channel to be imaged, it is important that the data set be small enough to fit in computer memory if at all possible. Separating the IFs into separate files will not interfere with the imaging and will help with the data set size problem: <source lang="text"> > DEFAULT UVCOP ; INP to reset all adverbs and choose the task. > INDI Tn; GETN Tm CR to select the calibrated target data set on disk T n and catalog number Tm. > DOWAIT 1 CR to have the task resume AIPS only after it has finished. > OUTSEQ 0 ; OUTDISK INDISK CR to avoid file name issues and select the output disk. > FORBIF=1TON;EIF=BIF;END CR to make separate files of each of the N IFs. > DOWAIT -1 CR to turn off waiting. </source> OSRO data sets may not need this operation and skipping the above step will simplify any continuum imaging that you may wish to do. Doing this UVCOP step on large RSRO data sets will be worth any extra trouble it may cause. Note that you could perform the separation into separate IFs before UVLSF which will speed up POSSM and UVLSF. However, the continuum output would then have to be assembled using VBGLU, which is why the steps above were shown in the present order. Spectral-line imaging is discussed in § 8.4 as well as throughout Chapter 5. With large numbers of spectral channels, you may wish to have IMAGR find appropriate Clean boxes for you. Set IM2PARM(1) through IM2PARM(6) cautiously. IM2PARM(7) controls whether the boxes of channel n are passed on to channel n + 1. The default does not pass the boxes along when autoboxing which is probably the correct decision. The end result of the imaging will be one image “cube” for each IF since each IF has to be imaged separately even with a multi-ID input data set. (If you set BIF = 1; EIF = 0 and try to image channel 103, you will actually image the average of channel 103 from each of the IFs.) To put the individual cubes together into one large cube, use MCUBE (§ 8.5.1). == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] Figure E.1: The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten. 96cb532ffb7370a93b52cb7ec82be30744818e54 32 31 2010-04-13T14:48:24Z Jmcmulli 2 /* Initial Editing */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. '''Note:''' This document is written as an appendix to the AIPS Cookbook; section numbers refer to that document: http://www.aips.nrao.edu/cook.html == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 CR to apply the delay calibration — very important. > SOLINT 0 CR to compute a bandpass solution for each scan on the bandpass calibrator. > CALSOUR ’bandpass cal’ CR to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. Remember these values — you will use them again. > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR to normalize the results only after the solution is found using the channels selected by ICHANSEL. > GO CR to make a bandpass (BP) table. </source> Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: <source lang="text"> > DEFAULT CALIB ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IN2DI n2; GET2N m2 CR to select the model image on disk n2 and catalog number m2. > DOCAL1;DOBAND3 CR to apply the delay and bandpass calibration — very important. > SOLINT 0 ; NMAPS 1 CR to compute a solution for each calibration scan and use the source model. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. </source> Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: <source lang="text"> > CLR2NAME ; NMAPS 0 CR to do no models. > CALSOUR ’other cal1’, ’other cal2’ CR to select the secondary calibrators by whatever names appear in your LISTR output. > GO CR to find the remaining complex gains. </source> Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: <source lang="text"> > TASK ’GETJY’ ; INP CR to set the task name without changing other adverbs. > SOURCE CALSOUR CR to select the secondary sources by the list of name you just used. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > INP CR to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. > GO CR to adjust the gains in the SN table and the fluxes in the SU (source) table. </source> Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: <source lang="text"> > DEFAULT CLCAL ; INP CR to clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > CALCODE ’*’ CR to select all calibration sources. > SNVER 2; INVERS SNVER CR to select your solution table from CALIB. Do not include the SN table from FRING a second time! > GO CR to apply SN table 2 to CL table 2, creating CL table 3. </source> Check the result using POSSM and/or VPLOT. == Target Source Data -- Edit and SPLIT == At this point, your calibration should be finished. You should now do an initial editing on the target sources, much like that done above for the calibration sources: <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all channels into one number. > CALCODE ’-CAL’ CR to do just target sources now. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> Again, remember to set it to flag all channels and IFs. You may have to select sub-windows and force the averaging to one times ∆ t to edit in detail, or perhaps the default time averaging will be beneficial. In general, the DISPLAY AMP V DIFF is a powerful way to catch bad amplitudes and phases. It will catch drop outs either as bright lines for strong sources or dark grey ones for weak sources. Since EVLA data sets tend to be large and unwieldy, it is recommended that you separate the data into the separate target sources, applying the current calibration and flagging once and for all. The imaging task IMAGR can do this on the fly, but, especially for observations of spectral-line sources, this is excessively expensive. <source lang="text"> > DEFAULT SPLIT ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > CALCODE ’-CAL’ CR to do just target sources now. > GO CR to write out separate calibrated data sets for each target source. </source> At present, EVLA data sets have no meaningful weights associated with the visibility data. There is a task new to 31DEC10 called REWAY which computes a robust rms over spectral channels within each IF and polarization. It can simply base the weights on these on a record-by-record, baseline-by-baseline basis or it can average the rmses in time, solve for antenna-based rmses, smooth those further in time, and only then apply them to the data. For these weights to be meaningful, the bandpass calibration must be applied and any RFI or other real spectral-line signal channels must be omitted from the rms computation. For the weights to be correctly calibrated, all amplitude calibration must also be applied. For these reasons, REWAY might well be used instead of SPLIT, running it one source at a time. Thus, <source lang="text"> > DEFAULT REWAY ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > SOURCE ’target1’ , ’ ’ CR to do one target source. > GO CR to write out a calibrated, weighted data set for the first target source. </source> Then, when that finishes <source lang="text"> > SOURCE ’target2’ , ’ ’ ; GO CR to do another target source. </source> It is not clear that this algorithm is optimal, but it certainly should be better than using all weights 1.0 throughout. == Spectral-line Imaging Hints == In many spectral-line observations you will now want to separate the continuum signal from the channel- dependent signals. This is discussed in some detail in §8.3. The larger number of channels from the EVLA does mean that continuum may be estimated with greater accuracy than when there were rather few channels which were both free of edge effects and spectral-line signal. The wider total bandwidth may, however, invalidate the assumption that the continuum signal at each visibility point can be represented by a polynomial of zero or first order. If there is a single dominant continuum source offset from the phase center, the assumption may be rendered valid by shifting the data with UVLSF to center the continuum source temporarily in order to subtract it. To examine this assumption and to determine which channels appear safe to use as “continuum” channels, use POSSM. <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array > BASELINE ANTEN CR and only them. > APARM 1 , 0 CR to display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels which will assist in determining channels that should be omitted. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> Note also whether the continuum appears to be a linear function of channel. If so, then use UVLSF to fit the continuum signal, writing a continuum only and a spectral-line only data set: <source lang="text"> > DEFAULT UVLSF ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for fitting the continuum. For a multi-IF data set, you will need to select the channel ranges carefully by IF. > ORDER 1 CR to select fitting the continuum in real and imaginary parts with a first order polynomial in channel number. UVLSF offers orders up to four, but they are not for the faint at heart and will give bad results if there are large ranges of channels left out of the fit due to line signals. > DOOUTPUT 1 CR to have the continuum which was fit written as a separate data set. This may be used to image the continuum. > SHIFT ∆x,∆y CR to shift the phase center to the dominant continuum source temporarily for the fitting. > GO CR to run the task. </source> Imaging the continuum output may, in addition to any scientific value of the continuum image, provide additional flagging and even self-calibration information which may be applied to the line data. If UVLSF cannot be used, flag the channels at the edges and those with spectral signals using UVFLG. Construct a continuum image with IMAGR on this flagged, spectral-line data set. Note that you might want to reduce the size of the data set with time averaging (UVAVG) and/or channel averaging (SPLIT or AVSPC) before beginning the imaging. Imaging is discussed in detail in § 5.2 through § 5.3.6 and will not be discussed here. You may find that additional editing is needed and that iterative self-calibration is of use. Be sure to copy those flags (but not the edge and spectral-signal flags) and final SN table back to the line data set. Apply them with SPLIT and then subtract the final continuum model with UVSUB. It you have had to use the spectral-index options of IMAGR, you may do the proper subtraction including these options with OOSUB rather than UVSUB. Spectral-line imaging of EVLA data will resemble that for the old VLA except for the increased number of spectral channels and the consequent increase in the data set size. Since IMAGR must read the full data set to select the data for the next channel to be imaged, it is important that the data set be small enough to fit in computer memory if at all possible. Separating the IFs into separate files will not interfere with the imaging and will help with the data set size problem: <source lang="text"> > DEFAULT UVCOP ; INP to reset all adverbs and choose the task. > INDI Tn; GETN Tm CR to select the calibrated target data set on disk T n and catalog number Tm. > DOWAIT 1 CR to have the task resume AIPS only after it has finished. > OUTSEQ 0 ; OUTDISK INDISK CR to avoid file name issues and select the output disk. > FORBIF=1TON;EIF=BIF;END CR to make separate files of each of the N IFs. > DOWAIT -1 CR to turn off waiting. </source> OSRO data sets may not need this operation and skipping the above step will simplify any continuum imaging that you may wish to do. Doing this UVCOP step on large RSRO data sets will be worth any extra trouble it may cause. Note that you could perform the separation into separate IFs before UVLSF which will speed up POSSM and UVLSF. However, the continuum output would then have to be assembled using VBGLU, which is why the steps above were shown in the present order. Spectral-line imaging is discussed in § 8.4 as well as throughout Chapter 5. With large numbers of spectral channels, you may wish to have IMAGR find appropriate Clean boxes for you. Set IM2PARM(1) through IM2PARM(6) cautiously. IM2PARM(7) controls whether the boxes of channel n are passed on to channel n + 1. The default does not pass the boxes along when autoboxing which is probably the correct decision. The end result of the imaging will be one image “cube” for each IF since each IF has to be imaged separately even with a multi-ID input data set. (If you set BIF = 1; EIF = 0 and try to image channel 103, you will actually image the average of channel 103 from each of the IFs.) To put the individual cubes together into one large cube, use MCUBE (§ 8.5.1). == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] Figure E.1: The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten. b35e223a3f901436124307d8597722224f824653 33 32 2010-04-13T14:50:59Z Jmcmulli 2 /* Basic Calibration */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. '''Note:''' This document is written as an appendix to the AIPS Cookbook; section numbers refer to that document: http://www.aips.nrao.edu/cook.html == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 CR to apply the delay calibration — very important. > SOLINT 0 CR to compute a bandpass solution for each scan on the bandpass calibrator. > CALSOUR ’bandpass cal’ CR to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. Remember these values — you will use them again. > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR to normalize the results only after the solution is found using the channels selected by ICHANSEL. > GO CR to make a bandpass (BP) table. </source> Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: <source lang="text"> > DEFAULT CALIB ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IN2DI n2; GET2N m2 CR to select the model image on disk n2 and catalog number m2. > DOCAL1;DOBAND3 CR to apply the delay and bandpass calibration — very important. > SOLINT 0 ; NMAPS 1 CR to compute a solution for each calibration scan and use the source model. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. </source> Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: <source lang="text"> > CLR2NAME ; NMAPS 0 CR to do no models. > CALSOUR ’other cal1’, ’other cal2’ CR to select the secondary calibrators by whatever names appear in your LISTR output. > GO CR to find the remaining complex gains. </source> Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: <source lang="text"> > TASK ’GETJY’ ; INP CR to set the task name without changing other adverbs. > SOURCE CALSOUR CR to select the secondary sources by the list of name you just used. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > INP CR to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. > GO CR to adjust the gains in the SN table and the fluxes in the SU (source) table. </source> Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: <source lang="text"> > DEFAULT CLCAL ; INP CR to clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > CALCODE ’*’ CR to select all calibration sources. > SNVER 2; INVERS SNVER CR to select your solution table from CALIB. Do not include the SN table from FRING a second time! > GO CR to apply SN table 2 to CL table 2, creating CL table 3. </source> Check the result using POSSM and/or VPLOT. == Target Source Data -- Edit and SPLIT == At this point, your calibration should be finished. You should now do an initial editing on the target sources, much like that done above for the calibration sources: <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all channels into one number. > CALCODE ’-CAL’ CR to do just target sources now. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> Again, remember to set it to flag all channels and IFs. You may have to select sub-windows and force the averaging to one times ∆ t to edit in detail, or perhaps the default time averaging will be beneficial. In general, the DISPLAY AMP V DIFF is a powerful way to catch bad amplitudes and phases. It will catch drop outs either as bright lines for strong sources or dark grey ones for weak sources. Since EVLA data sets tend to be large and unwieldy, it is recommended that you separate the data into the separate target sources, applying the current calibration and flagging once and for all. The imaging task IMAGR can do this on the fly, but, especially for observations of spectral-line sources, this is excessively expensive. <source lang="text"> > DEFAULT SPLIT ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > CALCODE ’-CAL’ CR to do just target sources now. > GO CR to write out separate calibrated data sets for each target source. </source> At present, EVLA data sets have no meaningful weights associated with the visibility data. There is a task new to 31DEC10 called REWAY which computes a robust rms over spectral channels within each IF and polarization. It can simply base the weights on these on a record-by-record, baseline-by-baseline basis or it can average the rmses in time, solve for antenna-based rmses, smooth those further in time, and only then apply them to the data. For these weights to be meaningful, the bandpass calibration must be applied and any RFI or other real spectral-line signal channels must be omitted from the rms computation. For the weights to be correctly calibrated, all amplitude calibration must also be applied. For these reasons, REWAY might well be used instead of SPLIT, running it one source at a time. Thus, <source lang="text"> > DEFAULT REWAY ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > SOURCE ’target1’ , ’ ’ CR to do one target source. > GO CR to write out a calibrated, weighted data set for the first target source. </source> Then, when that finishes <source lang="text"> > SOURCE ’target2’ , ’ ’ ; GO CR to do another target source. </source> It is not clear that this algorithm is optimal, but it certainly should be better than using all weights 1.0 throughout. == Spectral-line Imaging Hints == In many spectral-line observations you will now want to separate the continuum signal from the channel- dependent signals. This is discussed in some detail in §8.3. The larger number of channels from the EVLA does mean that continuum may be estimated with greater accuracy than when there were rather few channels which were both free of edge effects and spectral-line signal. The wider total bandwidth may, however, invalidate the assumption that the continuum signal at each visibility point can be represented by a polynomial of zero or first order. If there is a single dominant continuum source offset from the phase center, the assumption may be rendered valid by shifting the data with UVLSF to center the continuum source temporarily in order to subtract it. To examine this assumption and to determine which channels appear safe to use as “continuum” channels, use POSSM. <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array > BASELINE ANTEN CR and only them. > APARM 1 , 0 CR to display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels which will assist in determining channels that should be omitted. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> Note also whether the continuum appears to be a linear function of channel. If so, then use UVLSF to fit the continuum signal, writing a continuum only and a spectral-line only data set: <source lang="text"> > DEFAULT UVLSF ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for fitting the continuum. For a multi-IF data set, you will need to select the channel ranges carefully by IF. > ORDER 1 CR to select fitting the continuum in real and imaginary parts with a first order polynomial in channel number. UVLSF offers orders up to four, but they are not for the faint at heart and will give bad results if there are large ranges of channels left out of the fit due to line signals. > DOOUTPUT 1 CR to have the continuum which was fit written as a separate data set. This may be used to image the continuum. > SHIFT ∆x,∆y CR to shift the phase center to the dominant continuum source temporarily for the fitting. > GO CR to run the task. </source> Imaging the continuum output may, in addition to any scientific value of the continuum image, provide additional flagging and even self-calibration information which may be applied to the line data. If UVLSF cannot be used, flag the channels at the edges and those with spectral signals using UVFLG. Construct a continuum image with IMAGR on this flagged, spectral-line data set. Note that you might want to reduce the size of the data set with time averaging (UVAVG) and/or channel averaging (SPLIT or AVSPC) before beginning the imaging. Imaging is discussed in detail in § 5.2 through § 5.3.6 and will not be discussed here. You may find that additional editing is needed and that iterative self-calibration is of use. Be sure to copy those flags (but not the edge and spectral-signal flags) and final SN table back to the line data set. Apply them with SPLIT and then subtract the final continuum model with UVSUB. It you have had to use the spectral-index options of IMAGR, you may do the proper subtraction including these options with OOSUB rather than UVSUB. Spectral-line imaging of EVLA data will resemble that for the old VLA except for the increased number of spectral channels and the consequent increase in the data set size. Since IMAGR must read the full data set to select the data for the next channel to be imaged, it is important that the data set be small enough to fit in computer memory if at all possible. Separating the IFs into separate files will not interfere with the imaging and will help with the data set size problem: <source lang="text"> > DEFAULT UVCOP ; INP to reset all adverbs and choose the task. > INDI Tn; GETN Tm CR to select the calibrated target data set on disk T n and catalog number Tm. > DOWAIT 1 CR to have the task resume AIPS only after it has finished. > OUTSEQ 0 ; OUTDISK INDISK CR to avoid file name issues and select the output disk. > FORBIF=1TON;EIF=BIF;END CR to make separate files of each of the N IFs. > DOWAIT -1 CR to turn off waiting. </source> OSRO data sets may not need this operation and skipping the above step will simplify any continuum imaging that you may wish to do. Doing this UVCOP step on large RSRO data sets will be worth any extra trouble it may cause. Note that you could perform the separation into separate IFs before UVLSF which will speed up POSSM and UVLSF. However, the continuum output would then have to be assembled using VBGLU, which is why the steps above were shown in the present order. Spectral-line imaging is discussed in § 8.4 as well as throughout Chapter 5. With large numbers of spectral channels, you may wish to have IMAGR find appropriate Clean boxes for you. Set IM2PARM(1) through IM2PARM(6) cautiously. IM2PARM(7) controls whether the boxes of channel n are passed on to channel n + 1. The default does not pass the boxes along when autoboxing which is probably the correct decision. The end result of the imaging will be one image “cube” for each IF since each IF has to be imaged separately even with a multi-ID input data set. (If you set BIF = 1; EIF = 0 and try to image channel 103, you will actually image the average of channel 103 from each of the IFs.) To put the individual cubes together into one large cube, use MCUBE (§ 8.5.1). == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] Figure E.1: The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten. 93491953669bf0976c1f2e4b3442f4802de7ecb6 34 33 2010-04-13T14:53:03Z Jmcmulli 2 /* Target Source Data -- Edit and SPLIT */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. '''Note:''' This document is written as an appendix to the AIPS Cookbook; section numbers refer to that document: http://www.aips.nrao.edu/cook.html == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 CR to apply the delay calibration — very important. > SOLINT 0 CR to compute a bandpass solution for each scan on the bandpass calibrator. > CALSOUR ’bandpass cal’ CR to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. Remember these values — you will use them again. > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR to normalize the results only after the solution is found using the channels selected by ICHANSEL. > GO CR to make a bandpass (BP) table. </source> Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: <source lang="text"> > DEFAULT CALIB ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IN2DI n2; GET2N m2 CR to select the model image on disk n2 and catalog number m2. > DOCAL1;DOBAND3 CR to apply the delay and bandpass calibration — very important. > SOLINT 0 ; NMAPS 1 CR to compute a solution for each calibration scan and use the source model. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. </source> Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: <source lang="text"> > CLR2NAME ; NMAPS 0 CR to do no models. > CALSOUR ’other cal1’, ’other cal2’ CR to select the secondary calibrators by whatever names appear in your LISTR output. > GO CR to find the remaining complex gains. </source> Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: <source lang="text"> > TASK ’GETJY’ ; INP CR to set the task name without changing other adverbs. > SOURCE CALSOUR CR to select the secondary sources by the list of name you just used. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > INP CR to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. > GO CR to adjust the gains in the SN table and the fluxes in the SU (source) table. </source> Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: <source lang="text"> > DEFAULT CLCAL ; INP CR to clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > CALCODE ’*’ CR to select all calibration sources. > SNVER 2; INVERS SNVER CR to select your solution table from CALIB. Do not include the SN table from FRING a second time! > GO CR to apply SN table 2 to CL table 2, creating CL table 3. </source> Check the result using POSSM and/or VPLOT. == Target Source Data -- Edit and SPLIT == At this point, your calibration should be finished. You should now do an initial editing on the target sources, much like that done above for the calibration sources: <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all channels into one number. > CALCODE ’-CAL’ CR to do just target sources now. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> Again, remember to set it to flag all channels and IFs. You may have to select sub-windows and force the averaging to one times ∆ t to edit in detail, or perhaps the default time averaging will be beneficial. In general, the DISPLAY AMP V DIFF is a powerful way to catch bad amplitudes and phases. It will catch drop outs either as bright lines for strong sources or dark grey ones for weak sources. Since EVLA data sets tend to be large and unwieldy, it is recommended that you separate the data into the separate target sources, applying the current calibration and flagging once and for all. The imaging task IMAGR can do this on the fly, but, especially for observations of spectral-line sources, this is excessively expensive. <source lang="text"> > DEFAULT SPLIT ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > CALCODE ’-CAL’ CR to do just target sources now. > GO CR to write out separate calibrated data sets for each target source. </source> At present, EVLA data sets have no meaningful weights associated with the visibility data. There is a task new to 31DEC10 called REWAY which computes a robust rms over spectral channels within each IF and polarization. It can simply base the weights on these on a record-by-record, baseline-by-baseline basis or it can average the rmses in time, solve for antenna-based rmses, smooth those further in time, and only then apply them to the data. For these weights to be meaningful, the bandpass calibration must be applied and any RFI or other real spectral-line signal channels must be omitted from the rms computation. For the weights to be correctly calibrated, all amplitude calibration must also be applied. For these reasons, REWAY might well be used instead of SPLIT, running it one source at a time. Thus, <source lang="text"> > DEFAULT REWAY ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > SOURCE ’target1’ , ’ ’ CR to do one target source. > GO CR to write out a calibrated, weighted data set for the 1st target source. </source> Then, when that finishes <source lang="text"> > SOURCE ’target2’ , ’ ’ ; GO CR to do another target source. </source> It is not clear that this algorithm is optimal, but it certainly should be better than using all weights 1.0 throughout. == Spectral-line Imaging Hints == In many spectral-line observations you will now want to separate the continuum signal from the channel- dependent signals. This is discussed in some detail in §8.3. The larger number of channels from the EVLA does mean that continuum may be estimated with greater accuracy than when there were rather few channels which were both free of edge effects and spectral-line signal. The wider total bandwidth may, however, invalidate the assumption that the continuum signal at each visibility point can be represented by a polynomial of zero or first order. If there is a single dominant continuum source offset from the phase center, the assumption may be rendered valid by shifting the data with UVLSF to center the continuum source temporarily in order to subtract it. To examine this assumption and to determine which channels appear safe to use as “continuum” channels, use POSSM. <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array > BASELINE ANTEN CR and only them. > APARM 1 , 0 CR to display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels which will assist in determining channels that should be omitted. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> Note also whether the continuum appears to be a linear function of channel. If so, then use UVLSF to fit the continuum signal, writing a continuum only and a spectral-line only data set: <source lang="text"> > DEFAULT UVLSF ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for fitting the continuum. For a multi-IF data set, you will need to select the channel ranges carefully by IF. > ORDER 1 CR to select fitting the continuum in real and imaginary parts with a first order polynomial in channel number. UVLSF offers orders up to four, but they are not for the faint at heart and will give bad results if there are large ranges of channels left out of the fit due to line signals. > DOOUTPUT 1 CR to have the continuum which was fit written as a separate data set. This may be used to image the continuum. > SHIFT ∆x,∆y CR to shift the phase center to the dominant continuum source temporarily for the fitting. > GO CR to run the task. </source> Imaging the continuum output may, in addition to any scientific value of the continuum image, provide additional flagging and even self-calibration information which may be applied to the line data. If UVLSF cannot be used, flag the channels at the edges and those with spectral signals using UVFLG. Construct a continuum image with IMAGR on this flagged, spectral-line data set. Note that you might want to reduce the size of the data set with time averaging (UVAVG) and/or channel averaging (SPLIT or AVSPC) before beginning the imaging. Imaging is discussed in detail in § 5.2 through § 5.3.6 and will not be discussed here. You may find that additional editing is needed and that iterative self-calibration is of use. Be sure to copy those flags (but not the edge and spectral-signal flags) and final SN table back to the line data set. Apply them with SPLIT and then subtract the final continuum model with UVSUB. It you have had to use the spectral-index options of IMAGR, you may do the proper subtraction including these options with OOSUB rather than UVSUB. Spectral-line imaging of EVLA data will resemble that for the old VLA except for the increased number of spectral channels and the consequent increase in the data set size. Since IMAGR must read the full data set to select the data for the next channel to be imaged, it is important that the data set be small enough to fit in computer memory if at all possible. Separating the IFs into separate files will not interfere with the imaging and will help with the data set size problem: <source lang="text"> > DEFAULT UVCOP ; INP to reset all adverbs and choose the task. > INDI Tn; GETN Tm CR to select the calibrated target data set on disk T n and catalog number Tm. > DOWAIT 1 CR to have the task resume AIPS only after it has finished. > OUTSEQ 0 ; OUTDISK INDISK CR to avoid file name issues and select the output disk. > FORBIF=1TON;EIF=BIF;END CR to make separate files of each of the N IFs. > DOWAIT -1 CR to turn off waiting. </source> OSRO data sets may not need this operation and skipping the above step will simplify any continuum imaging that you may wish to do. Doing this UVCOP step on large RSRO data sets will be worth any extra trouble it may cause. Note that you could perform the separation into separate IFs before UVLSF which will speed up POSSM and UVLSF. However, the continuum output would then have to be assembled using VBGLU, which is why the steps above were shown in the present order. Spectral-line imaging is discussed in § 8.4 as well as throughout Chapter 5. With large numbers of spectral channels, you may wish to have IMAGR find appropriate Clean boxes for you. Set IM2PARM(1) through IM2PARM(6) cautiously. IM2PARM(7) controls whether the boxes of channel n are passed on to channel n + 1. The default does not pass the boxes along when autoboxing which is probably the correct decision. The end result of the imaging will be one image “cube” for each IF since each IF has to be imaged separately even with a multi-ID input data set. (If you set BIF = 1; EIF = 0 and try to image channel 103, you will actually image the average of channel 103 from each of the IFs.) To put the individual cubes together into one large cube, use MCUBE (§ 8.5.1). == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] Figure E.1: The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten. 73839ff6daf6ee704d9417e2bce7cd776da7a921 35 34 2010-04-13T14:54:15Z Jmcmulli 2 /* Spectral-line Imaging Hints */ wikitext text/x-wiki [[Category:Post-Processing]][[Category:Calibration]][[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. '''Note:''' This document is written as an appendix to the AIPS Cookbook; section numbers refer to that document: http://www.aips.nrao.edu/cook.html == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 CR to apply the delay calibration — very important. > SOLINT 0 CR to compute a bandpass solution for each scan on the bandpass calibrator. > CALSOUR ’bandpass cal’ CR to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. Remember these values — you will use them again. > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR to normalize the results only after the solution is found using the channels selected by ICHANSEL. > GO CR to make a bandpass (BP) table. </source> Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: <source lang="text"> > DEFAULT CALIB ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IN2DI n2; GET2N m2 CR to select the model image on disk n2 and catalog number m2. > DOCAL1;DOBAND3 CR to apply the delay and bandpass calibration — very important. > SOLINT 0 ; NMAPS 1 CR to compute a solution for each calibration scan and use the source model. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. </source> Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: <source lang="text"> > CLR2NAME ; NMAPS 0 CR to do no models. > CALSOUR ’other cal1’, ’other cal2’ CR to select the secondary calibrators by whatever names appear in your LISTR output. > GO CR to find the remaining complex gains. </source> Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: <source lang="text"> > TASK ’GETJY’ ; INP CR to set the task name without changing other adverbs. > SOURCE CALSOUR CR to select the secondary sources by the list of name you just used. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > INP CR to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. > GO CR to adjust the gains in the SN table and the fluxes in the SU (source) table. </source> Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: <source lang="text"> > DEFAULT CLCAL ; INP CR to clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > CALCODE ’*’ CR to select all calibration sources. > SNVER 2; INVERS SNVER CR to select your solution table from CALIB. Do not include the SN table from FRING a second time! > GO CR to apply SN table 2 to CL table 2, creating CL table 3. </source> Check the result using POSSM and/or VPLOT. == Target Source Data -- Edit and SPLIT == At this point, your calibration should be finished. You should now do an initial editing on the target sources, much like that done above for the calibration sources: <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all channels into one number. > CALCODE ’-CAL’ CR to do just target sources now. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> Again, remember to set it to flag all channels and IFs. You may have to select sub-windows and force the averaging to one times ∆ t to edit in detail, or perhaps the default time averaging will be beneficial. In general, the DISPLAY AMP V DIFF is a powerful way to catch bad amplitudes and phases. It will catch drop outs either as bright lines for strong sources or dark grey ones for weak sources. Since EVLA data sets tend to be large and unwieldy, it is recommended that you separate the data into the separate target sources, applying the current calibration and flagging once and for all. The imaging task IMAGR can do this on the fly, but, especially for observations of spectral-line sources, this is excessively expensive. <source lang="text"> > DEFAULT SPLIT ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > CALCODE ’-CAL’ CR to do just target sources now. > GO CR to write out separate calibrated data sets for each target source. </source> At present, EVLA data sets have no meaningful weights associated with the visibility data. There is a task new to 31DEC10 called REWAY which computes a robust rms over spectral channels within each IF and polarization. It can simply base the weights on these on a record-by-record, baseline-by-baseline basis or it can average the rmses in time, solve for antenna-based rmses, smooth those further in time, and only then apply them to the data. For these weights to be meaningful, the bandpass calibration must be applied and any RFI or other real spectral-line signal channels must be omitted from the rms computation. For the weights to be correctly calibrated, all amplitude calibration must also be applied. For these reasons, REWAY might well be used instead of SPLIT, running it one source at a time. Thus, <source lang="text"> > DEFAULT REWAY ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > SOURCE ’target1’ , ’ ’ CR to do one target source. > GO CR to write out a calibrated, weighted data set for the 1st target source. </source> Then, when that finishes <source lang="text"> > SOURCE ’target2’ , ’ ’ ; GO CR to do another target source. </source> It is not clear that this algorithm is optimal, but it certainly should be better than using all weights 1.0 throughout. == Spectral-line Imaging Hints == In many spectral-line observations you will now want to separate the continuum signal from the channel- dependent signals. This is discussed in some detail in §8.3. The larger number of channels from the EVLA does mean that continuum may be estimated with greater accuracy than when there were rather few channels which were both free of edge effects and spectral-line signal. The wider total bandwidth may, however, invalidate the assumption that the continuum signal at each visibility point can be represented by a polynomial of zero or first order. If there is a single dominant continuum source offset from the phase center, the assumption may be rendered valid by shifting the data with UVLSF to center the continuum source temporarily in order to subtract it. To examine this assumption and to determine which channels appear safe to use as “continuum” channels, use POSSM. <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array > BASELINE ANTEN CR and only them. > APARM 1 , 0 CR to display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels which will assist in determining channels that should be omitted. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> Note also whether the continuum appears to be a linear function of channel. If so, then use UVLSF to fit the continuum signal, writing a continuum only and a spectral-line only data set: <source lang="text"> > DEFAULT UVLSF ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for fitting the continuum. For a multi-IF data set, you will need to select the channel ranges carefully by IF. > ORDER 1 CR to select fitting the continuum in real and imaginary parts with a first order polynomial in channel number. UVLSF offers orders up to four, but they are not for the faint at heart and will give bad results if there are large ranges of channels left out of the fit due to line signals. > DOOUTPUT 1 CR to have the continuum which was fit written as a separate data set. This may be used to image the continuum. > SHIFT ∆x,∆y CR to shift the phase center to the dominant continuum source temporarily for the fitting. > GO CR to run the task. </source> Imaging the continuum output may, in addition to any scientific value of the continuum image, provide additional flagging and even self-calibration information which may be applied to the line data. If UVLSF cannot be used, flag the channels at the edges and those with spectral signals using UVFLG. Construct a continuum image with IMAGR on this flagged, spectral-line data set. Note that you might want to reduce the size of the data set with time averaging (UVAVG) and/or channel averaging (SPLIT or AVSPC) before beginning the imaging. Imaging is discussed in detail in § 5.2 through § 5.3.6 and will not be discussed here. You may find that additional editing is needed and that iterative self-calibration is of use. Be sure to copy those flags (but not the edge and spectral-signal flags) and final SN table back to the line data set. Apply them with SPLIT and then subtract the final continuum model with UVSUB. It you have had to use the spectral-index options of IMAGR, you may do the proper subtraction including these options with OOSUB rather than UVSUB. Spectral-line imaging of EVLA data will resemble that for the old VLA except for the increased number of spectral channels and the consequent increase in the data set size. Since IMAGR must read the full data set to select the data for the next channel to be imaged, it is important that the data set be small enough to fit in computer memory if at all possible. Separating the IFs into separate files will not interfere with the imaging and will help with the data set size problem: <source lang="text"> > DEFAULT UVCOP ; INP to reset all adverbs and choose the task. > INDI Tn; GETN Tm CR to select the calibrated target data set on disk T n and catalog number Tm. > DOWAIT 1 CR to have the task resume AIPS only after it has finished. > OUTSEQ 0 ; OUTDISK INDISK CR to avoid file name issues and select the output disk. > FORBIF=1TON;EIF=BIF;END CR to make separate files of each of the N IFs. > DOWAIT -1 CR to turn off waiting. </source> OSRO data sets may not need this operation and skipping the above step will simplify any continuum imaging that you may wish to do. Doing this UVCOP step on large RSRO data sets will be worth any extra trouble it may cause. Note that you could perform the separation into separate IFs before UVLSF which will speed up POSSM and UVLSF. However, the continuum output would then have to be assembled using VBGLU, which is why the steps above were shown in the present order. Spectral-line imaging is discussed in § 8.4 as well as throughout Chapter 5. With large numbers of spectral channels, you may wish to have IMAGR find appropriate Clean boxes for you. Set IM2PARM(1) through IM2PARM(6) cautiously. IM2PARM(7) controls whether the boxes of channel n are passed on to channel n + 1. The default does not pass the boxes along when autoboxing which is probably the correct decision. The end result of the imaging will be one image “cube” for each IF since each IF has to be imaged separately even with a multi-ID input data set. (If you set BIF = 1; EIF = 0 and try to image channel 103, you will actually image the average of channel 103 from each of the IFs.) To put the individual cubes together into one large cube, use MCUBE (§ 8.5.1). == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] Figure E.1: The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten. 5c4f560f034eade6bfeb213816528618a02a30d5 36 35 2010-04-13T14:57:34Z Jmcmulli 2 wikitext text/x-wiki [[Category:Post-Processing]] [[Category:Calibration]] [[Category:AIPS]] = Special Considerations for EVLA Data Calibration and Imaging in AIPS = == Overview == The old VLA with its once state of the art, but now dated, correlator and electronic has been turned off. The new electronics and correlator of the EVLA has been turned on and made available to users. For the time being, this availability will be “Shared Risk Observing” of two forms: “Open” (OSRO) with limited capabilities and “Resident” (RSRO) with potentially unlimited capabilities. OSRO data are expected to be similar to those of the VLA initially although with more spectral channels and bandwidth. OSRO data for continuum science will have, in each of 2 tunable spectral windows, 64 spectral channels in each of 4 polarization products covering ≤ 128 MHz. For line work, OSRO data will have one spectral window with 256 spectral channels in both RR and LL polarizations covering ≤ 128 MHz total bandwidth. These capabilities will increase gradually with time. RSRO data may have many thousands of spectral channels and, in time, up to 8 GHz of bandwidth per polarization. AIPS software will be important to both programs, although RSRO data are expected, by management, to be processed primarily in CASA. At this writing, the EVLA has already produced amazing scientific results, but with considerable difficulties which are expected to be corrected over time. Delays are currently not set accurately, causing the data analysis to begin with the VLBI task FRING needed to correct large slopes in phase across the bandpass. The flagging information known to the on-line system (telescope off source and the like) is now transferred to CASA and so may reach AIPS in the form of already flagged data. That on-line flagging information is not transferred to AIPS in the form of a flag table and neither software package yet receives useful flagging data, such as the antenna has no receiver, from the correlator. Substantial flagging effort is therefore still required. Also, system temperatures and gains are not transferred as yet. This means that the data weights are not meaningful and should not be adjusted by the amplitude calibration. That calibration will not be as good as it will be since corrections for differences in gains and sensitivity in the directions of the primary and secondary calibrators and the target sources cannot be made. The weather table is now available with the data so that reasonable opacities may be determined. However, the “over-the-top”, table which is used in determining antenna positions, the frequency offset table, used in managing Doppler tracking, and the CQ table, used to correct amplitudes for spectral averaging in the presence of non-zero delays, are not available. Although these issues should be corrected, quite possibly early in 2010, the following guide will not assume that they have been completed. Steps that can be omitted or simplified when they are will be described. This appendix is written with the assumption that the reader is moderately familiar with AIPS as described in the preceding chapters. It is also written with the assumption that you are using the 31DEC10 or later releases of the software. '''Note:''' This document is written as an appendix to the AIPS Cookbook; section numbers refer to that document: http://www.aips.nrao.edu/cook.html == Getting Your Data into AIPS == Your EVLA data are stored as an “ALMA Science Data Model” (ASDM) format file in “SDMBDF” (Science Data Model Binary Data Format) in the NRAO archive. They may be read out of the archive in that format, a CASA measurement set format, or in an AIPS-friendly uvfits format. This last is produced by the CASA uvfits writing software. Go the the web page: http://archive.cv.nrao.edu/ and select the Advanced Query Tool. Fill out enough of the form to describe your data and submit the query. If the data are not yet public, you will need the Locked Project Access Key which may be obtained from the NRAO data analysts. To avoid the need for this key, you may log in to my.nrao.edu after which it will know if you are entitled to access to locked projects. The query will return a list of the data sets which meet your specifications. On this form, enter your e-mail address, choose AIPS Friendly names (almost certainly does not work), AIPS FITS under the EVLA-WIDAR section and choose the desired spectral and time averaging. If the delays are not accurately known, spectral averaging can be damaging to the data amplitudes. However, the data are recorded at one-second intervals which is rather short, making the data voluminous. Judicious averaging can help with data set size and processing times without compromising the science. Choose the data set(s) you wish to receive and submit the request. You will be told an estimate of the output data set sizes and the amount of time you will need to wait for the format translation to occur. A 19 Gbyte SDMBDF file run as a test with no averaging was estimated to produce a UVFITS file of 30.26 Gbytes and to take 103 minutes to prepare for download. That time assumes that your download job is the only one being performed. If your download fails, you will probably be told erroneously by e-mail that it worked. The output file will however be missing or incomplete. Try again before contacting NRAO for help. Unlocked files will be downloaded to the NRAO public ftp site: ftp://ftp.aoc.nrao.edu/e2earchive/ and you may then use ftp to copy the file to your computer. Locked files will go to a protected ftp site and you must use ftp to download those, even within NRAO. The instructions for downloading will be e-mailed to you. Be sure to specify binary for the copy. If you are located in the AOC in Socorro, you may set an environment variable to the archive location, e.g., <source lang="bash"> export E2E=/home/acorn2/ftp/pub/e2earchive CR for bash shells setenv E2E /home/acorn2/ftp/pub/e2earchive CR for C shells such as tcsh </source> and simply unlocked data files directly from the public download area. Note that the file will be deleted automatically after 48 hours in both public and protected data areas. The data file may be read from disk into AIPS using UVLOD or FITLD, using: <source lang='text'> > DEFAULT ’UVLOD’ ; INP CR to initialize and review the inputs needed. > DATAIN ’E2E:filename’ CR where filename is the disk file name in logical area E2E; (see § 3.10.3). > DOUVCOMP FALSE CR to write visibilities in uncompressed format. There are no weights at present, so there is no loss of information in compressed format, but the conversion from compressed format costs more than reading the larger data files. > OUTNA ’myname’ CR to set the AIPS name. > OUTCL ’ ’ CR to take default (UVDATA) class. > OUTSEQ 0 CR to take next higher sequence #. > OUTDI 3 CR to write the data to disk 3 (one with enough space). > INP CR to review the inputs. > GO CR to run the program when you’re satisfied with inputs. </source> Watch the messages from UVLOD to see where your data set goes and whether the task ran properly. When it is finished, check the output header: <source lang="text"> > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IMHEAD CR to examine the header. </source> Note that the header does not show the usual complement of AIPS extension files. CASA translates the on-line data into its internal format and then writes the uvfits file read by AIPS. Since CASA does not have files comparable to AIPS index and CL tables, it does not provide them. To build index and calibration tables, use; <source lang="text"> > TASK ’INDXR’ ; INP CR to select the task and review its inputs. > INFILE ’ ’;PRTLEV=0 CR to be sure not to use an input text file and to avoid excess messages. > CPARM=0,0,1/2 CR to make a CL table 1 with a 30-second interval. > BPARM τ , 0 CR to take default VLA gains and a zenith opacity of τ . Set τ = −1 for no opacity correction. You may set τ = 0, which is now recommended, to get new default opacities. These are based on a detailed model predicting the opacity at any frequency from that at 22 GHz. The combination of weather and seasonal model long used by FILLM and INDXR is now used solely to estimate the 22 GHz opacity. > GO CR to run the task after checking the inputs. </source> It is a good idea to list the structure of your data set and your antenna locations on the printer and to keep those listings next to your work station for reference: <source lang="text"> > DEFAULT LISTR ; INP CR to initialize the LISTR inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > OPTYPE ’SCAN’ ; DOCRT -1 CR to choose a scan listing on the printer. > GO ; GO PRTAN CR to print the scan listing and the antenna file contents. </source> Read these with care. There have sometimes been problems with antenna identifications, with the order of the IF frequencies, and even with identification of sources by scan. Task SUFIX may be used to correct the last problem and, if desired, FLOPM may be used to reverse the frequency order. You may have to use SETJY to change the CALCODE of some sources if your calibration sources have a blank calibrator code or your target sources have a non-blank calibrator code. == Initial Calibration -- FRING == We have had difficulty setting all of the delays in the EVLA to values which are sufficiently accurate. If the delay is not set correctly, the interferometer phase will vary linearly with frequency, potentially wrapping through several turns of phase within a single spectral window (“IF band”). We hope that bad delays will not arise in future, allowing you to skip this section. But this is a problem familiar to VLBI users and AIPS has a well-tested method to correct the problem. Using your LISTR output, select a time range of about one minute toward the end of a scan on a strong point-source calibrator, usually your bandpass calibrator. Then <source lang='text'> > DEFAULT FRING ; INP CR to initialize the FRING inputs and review them. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > TIMERANG db,hb,mb,sb,de,he,me,se CR to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. > SOLINT 1.05 CR to set the averaging interval in minutes slightly longer than the data interval selected. > DPARM(9) = 1 CR to fit only delay, not rate. > INP CR to check the voluminous inputs. > GO to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. </source> to specify the beginning day, hour, minute, and second and ending day, hour, minute, and second (wrt REFDATE) of the data to be included. Too much data will cause trouble. to set the averaging interval in minutes slightly longer than the data interval selected. to fit only delay, not rate. to check the voluminous inputs. to run the task, writing SN table 1 with delays for each antenna, IF, and polarization. The different IFs in current EVLA data sets may come from different basebands and therefore have different residual delays. The option APARM(5)=1 to force all IFs to have the same delay solution is therefore no longer appropriate. This SN table will need to be applied to the main CL table created by INDXR. <source lang="text"> > TASK ’CLCAL’ ; INP CR to look at the necessary inputs. > TIMERANG 0 CR to reset the time range. > GAINUSE 0 ; GAINVER 0 CR to select the highest CL table as input and write one higher as output (version 1 and 2, resp. in this case). > SNVER 1 ; INVER 1 CR to use only the SN table just created. > INP CR to review the inputs. > GO CR to make an update calibration table. </source> Be sure to apply this (or higher) CL table with DOCALIB 1 in all later steps. == Initial Editing == There will be data validity information prepared both by the on-line control software and by the WIDAR correlator and this information will in time be available as an initial flag table. The tasks above will have applied this table for you by default since FLAGVER 0. On-line flags may already have caused data to be flagged within your data set (but not deleted) by CASA. Unfortunately, at this writing, no flag table is present and, even when it does appear, it may not be fully reliable initially. Thus, we need to look at the data to flag out whatever remains of the time off source not flagged in CASA using on-line flagging information. There have also been drop outs in which the visibility is pure zero, typically for all channels and IFs and a single integration. The drop outs should now be handled by UVLOD and FITLD. Note, however, that CASA and FITLD pass along all data samples, including those that are fully flagged. This makes the data set rather larger than one might wish. Use UVCOP to remove all fully flagged data samples. Before doing this, use TVFLG to look for any more data samples that might need to be flagged fully. Check especially samples at the beginnings and ends of scans. Try <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND -1 CR to apply the delay calibration. If a bandpass has been determined, use DOBAND 3 or 1 to apply it. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all the channels into one number. > CALCODE ’*’ CR to do just calibrators for the moment. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> The default smoothing time shown in the display will probably be some multiple of ∆ t. Select sub-windows and change the smoothing time to one times the basic interval in order to edit in detail. Remember to change the initial setup so that the flags apply to all channels and all IFs. See § 4.4.3 for more information. We note here that some users feel that the data need to be inspected more carefully than with just an average of most of the channels. POSSM (below) may be of use to find RFI. Avoiding the worst of that, you may still wish to run TVFLG to look at the average of a few channels at a time. Use NCHAV and CHINC appropriately. Task SPFLG (§ 10.2.2) is the ultimate weapon when looking for channel-dependent difficulties, but is onerous when there are many baselines. == Basic Calibration == For both continuum and line observations, we must begin by determining which spectral channels are reliable and which are affected by the inevitable loss of signal-to-noise at band edges or are degraded by radio- frequency interference (RFI). Use POSSM to display spectra from the shorter baselines on the TV: <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > SOURCE ’bandpass cal’ CR to select the strong bandpass calibrator. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array or the maintenance areas. > BASELINE ANTEN CR and only them. > DOCAL 1;APARM 1,0 CR to apply the FRING solutions and display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> If there is no RFI, then you may be able to use the same channel range for all IFs. If the RFI is particularly pernicious, you may have to edit it out of your data before continuing. Task FLGIT (§ 8.1) attempts to flag RFI that is both channel- and time-dependent in a non-interactive fashion. SPFLG (§ 10.2.2) is labor and time intensive but would be the most reliable method to deal with the problem. The basic EVLA calibration is much like that described in detail in Chapter 4 except that bandpass calibration is now required rather than merely recommended. Having chosen those channels which may be reliably used to normalize the bandpass functions, <source lang="text"> > DEFAULT BPASS ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 CR to apply the delay calibration — very important. > SOLINT 0 CR to compute a bandpass solution for each scan on the bandpass calibrator. > CALSOUR ’bandpass cal’ CR to select the strong bandpass calibrator. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. Remember these values — you will use them again. > BPASSPRM(5) 1 ; BPASSPRM(10) 3 CR to normalize the results only after the solution is found using the channels selected by ICHANSEL. > GO CR to make a bandpass (BP) table. </source> Do not use spectral smoothing at this point unless you want to use the same smoothing forever after. Apply the flag table. Consider correcting the bandpass function for the spectral index of bandpass cal if it is known — the EVLA bandwidths are large enough that this may matter. A model for the calibrator may be used; see § 4.3.3.1. You now need to run SETJY with OPTYPE ’CALC’ and SOURCES set to point at your primary flux calibration sources. You should load the models for these sources that apply to your data with CALRD; see §4.3.3.1. Then run CALIB with the model once for each primary flux calibrator: <source lang="text"> > DEFAULT CALIB ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > IN2DI n2; GET2N m2 CR to select the model image on disk n2 and catalog number m2. > DOCAL1;DOBAND3 CR to apply the delay and bandpass calibration — very important. > SOLINT 0 ; NMAPS 1 CR to compute a solution for each calibration scan and use the source model. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for averaging in each IF. These must be the same values that you used in BPASS. > SNVER 2 CR to put all CALIB solutions in solution table 2. > GO CR to find the complex gains for the flux calibrator. </source> Read the output closely. If solutions fail, examine your data closely for bad things. The primary flux calibrator should work without failure. After you have done each primary flux calibrator for which you have models, run CALIB on the remaining calibration sources: <source lang="text"> > CLR2NAME ; NMAPS 0 CR to do no models. > CALSOUR ’other cal1’, ’other cal2’ CR to select the secondary calibrators by whatever names appear in your LISTR output. > GO CR to find the remaining complex gains. </source> Again, examine the output messages closely. There may be a few failures but there should not be many in a good data set. The RUN file procedure VLACALIB (see § 4.3.3.1) may be used but it does not offer the ICHANSEL option which may be required by your data. It also does a scalar averaging for the amplitudes. In 31DEC10, this averaging was changed to be a vector average of the spectral channels followed by a scalar average over time. Scalar averaging suffers from Ricean bias in the amplitudes and so should be used only when the calibration source is very strong or when the atmospheric phases are very unstable. At this point it is necessary to calibrate the fluxes of the secondary calibration sources using your SN table: <source lang="text"> > TASK ’GETJY’ ; INP CR to set the task name without changing other adverbs. > SOURCE CALSOUR CR to select the secondary sources by the list of name you just used. > CALSOUR ’flux cal’ CR to select the primary flux calibrator by whatever form of its name appears in your LISTR output. > INP CR to check the inputs closely; remember to do all times, IFs, etc. with SNVER 2. > GO CR to adjust the gains in the SN table and the fluxes in the SU (source) table. </source> Look at the messages with care — the fluxes in the various IFs should be consistent and the error bars should be reasonably small (< 10% at high frequencies, smaller at low frequencies). If not, look at your SN table with SNPLT to see if there are bad solutions. If there are, delete SN table 2, do more flagging with TVFLG or SPFLG, and repeat the process. Finally, apply the gain solutions to your calibration table: <source lang="text"> > DEFAULT CLCAL ; INP CR to clear the adverbs. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > CALCODE ’*’ CR to select all calibration sources. > SNVER 2; INVERS SNVER CR to select your solution table from CALIB. Do not include the SN table from FRING a second time! > GO CR to apply SN table 2 to CL table 2, creating CL table 3. </source> Check the result using POSSM and/or VPLOT. == Target Source Data -- Edit and SPLIT == At this point, your calibration should be finished. You should now do an initial editing on the target sources, much like that done above for the calibration sources: <source lang="text"> > DEFAULT TVFLG ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > BCHAN c1 ; ECHAN c2 CR to average across a range of channels — not as flexible as ICHANSEL but probably okay here. > BIF j ; EIF BIF CR to edit one IF only, which will suffice for problems that are not IF dependent, such as drop outs, antenna not on source, etc. Choose an IF that is reasonably free of RFI. > NCHAV ECHAN-BCHAN+1 CR to average all channels into one number. > CALCODE ’-CAL’ CR to do just target sources now. > DPARM(6) ∆ t CR to do no time averaging in the work file set ∆ t to the data interval in seconds. > GO CR to start the task. </source> Again, remember to set it to flag all channels and IFs. You may have to select sub-windows and force the averaging to one times ∆ t to edit in detail, or perhaps the default time averaging will be beneficial. In general, the DISPLAY AMP V DIFF is a powerful way to catch bad amplitudes and phases. It will catch drop outs either as bright lines for strong sources or dark grey ones for weak sources. Since EVLA data sets tend to be large and unwieldy, it is recommended that you separate the data into the separate target sources, applying the current calibration and flagging once and for all. The imaging task IMAGR can do this on the fly, but, especially for observations of spectral-line sources, this is excessively expensive. <source lang="text"> > DEFAULT SPLIT ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > CALCODE ’-CAL’ CR to do just target sources now. > GO CR to write out separate calibrated data sets for each target source. </source> At present, EVLA data sets have no meaningful weights associated with the visibility data. There is a task new to 31DEC10 called REWAY which computes a robust rms over spectral channels within each IF and polarization. It can simply base the weights on these on a record-by-record, baseline-by-baseline basis or it can average the rmses in time, solve for antenna-based rmses, smooth those further in time, and only then apply them to the data. For these weights to be meaningful, the bandpass calibration must be applied and any RFI or other real spectral-line signal channels must be omitted from the rms computation. For the weights to be correctly calibrated, all amplitude calibration must also be applied. For these reasons, REWAY might well be used instead of SPLIT, running it one source at a time. Thus, <source lang="text"> > DEFAULT REWAY ; INP to reset all adverbs and choose the task. > INDI n; GETN m CR to select the data set on disk n and catalog number m. > DOCAL 1 ; DOBAND 3 CR to apply the delay, complex gain, and bandpass calibration. > SOURCE ’target1’ , ’ ’ CR to do one target source. > GO CR to write out a calibrated, weighted data set for the 1st target source. </source> Then, when that finishes <source lang="text"> > SOURCE ’target2’ , ’ ’ ; GO CR to do another target source. </source> It is not clear that this algorithm is optimal, but it certainly should be better than using all weights 1.0 throughout. == Spectral-line Imaging Hints == In many spectral-line observations you will now want to separate the continuum signal from the channel- dependent signals. This is discussed in some detail in §8.3. The larger number of channels from the EVLA does mean that continuum may be estimated with greater accuracy than when there were rather few channels which were both free of edge effects and spectral-line signal. The wider total bandwidth may, however, invalidate the assumption that the continuum signal at each visibility point can be represented by a polynomial of zero or first order. If there is a single dominant continuum source offset from the phase center, the assumption may be rendered valid by shifting the data with UVLSF to center the continuum source temporarily in order to subtract it. To examine this assumption and to determine which channels appear safe to use as “continuum” channels, use POSSM. <source lang="text"> > DEFAULT POSSM ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > DOTV 1;NPLOTS 1 CR to plot only on the TV, one baseline at a time. > ANTEN n1,n2,n3,n4 CR to select the antennas nearest the center of the array > BASELINE ANTEN CR and only them. > APARM 1 , 0 CR to display vector averaged spectra. Scalar averaged spectra will turn up at the edges reflecting the decreased signal to noise in the outer channels which will assist in determining channels that should be omitted. > BIF j ; EIF BIF CR to plot one IF at a time. > GO CR to run the task. Make notes of the desirable channels IF by IF. </source> Note also whether the continuum appears to be a linear function of channel. If so, then use UVLSF to fit the continuum signal, writing a continuum only and a spectral-line only data set: <source lang="text"> > DEFAULT UVLSF ; INP to set the task name and clear the adverbs. > INDI Tn; GETN Tm CR to select the calibrated line data set on disk Tn and catalog number Tm. > ICHANSEL c11,c12,1,if1,c21,c22,1,if2,c31,c32,1,if3,... CR to select the range(s) of channels which are reliable for fitting the continuum. For a multi-IF data set, you will need to select the channel ranges carefully by IF. > ORDER 1 CR to select fitting the continuum in real and imaginary parts with a first order polynomial in channel number. UVLSF offers orders up to four, but they are not for the faint at heart and will give bad results if there are large ranges of channels left out of the fit due to line signals. > DOOUTPUT 1 CR to have the continuum which was fit written as a separate data set. This may be used to image the continuum. > SHIFT ∆x,∆y CR to shift the phase center to the dominant continuum source temporarily for the fitting. > GO CR to run the task. </source> Imaging the continuum output may, in addition to any scientific value of the continuum image, provide additional flagging and even self-calibration information which may be applied to the line data. If UVLSF cannot be used, flag the channels at the edges and those with spectral signals using UVFLG. Construct a continuum image with IMAGR on this flagged, spectral-line data set. Note that you might want to reduce the size of the data set with time averaging (UVAVG) and/or channel averaging (SPLIT or AVSPC) before beginning the imaging. Imaging is discussed in detail in § 5.2 through § 5.3.6 and will not be discussed here. You may find that additional editing is needed and that iterative self-calibration is of use. Be sure to copy those flags (but not the edge and spectral-signal flags) and final SN table back to the line data set. Apply them with SPLIT and then subtract the final continuum model with UVSUB. It you have had to use the spectral-index options of IMAGR, you may do the proper subtraction including these options with OOSUB rather than UVSUB. Spectral-line imaging of EVLA data will resemble that for the old VLA except for the increased number of spectral channels and the consequent increase in the data set size. Since IMAGR must read the full data set to select the data for the next channel to be imaged, it is important that the data set be small enough to fit in computer memory if at all possible. Separating the IFs into separate files will not interfere with the imaging and will help with the data set size problem: <source lang="text"> > DEFAULT UVCOP ; INP to reset all adverbs and choose the task. > INDI Tn; GETN Tm CR to select the calibrated target data set on disk T n and catalog number Tm. > DOWAIT 1 CR to have the task resume AIPS only after it has finished. > OUTSEQ 0 ; OUTDISK INDISK CR to avoid file name issues and select the output disk. > FORBIF=1TON;EIF=BIF;END CR to make separate files of each of the N IFs. > DOWAIT -1 CR to turn off waiting. </source> OSRO data sets may not need this operation and skipping the above step will simplify any continuum imaging that you may wish to do. Doing this UVCOP step on large RSRO data sets will be worth any extra trouble it may cause. Note that you could perform the separation into separate IFs before UVLSF which will speed up POSSM and UVLSF. However, the continuum output would then have to be assembled using VBGLU, which is why the steps above were shown in the present order. Spectral-line imaging is discussed in § 8.4 as well as throughout Chapter 5. With large numbers of spectral channels, you may wish to have IMAGR find appropriate Clean boxes for you. Set IM2PARM(1) through IM2PARM(6) cautiously. IM2PARM(7) controls whether the boxes of channel n are passed on to channel n + 1. The default does not pass the boxes along when autoboxing which is probably the correct decision. The end result of the imaging will be one image “cube” for each IF since each IF has to be imaged separately even with a multi-ID input data set. (If you set BIF = 1; EIF = 0 and try to image channel 103, you will actually image the average of channel 103 from each of the IFs.) To put the individual cubes together into one large cube, use MCUBE (§ 8.5.1). == Continuum Imaging Hints == The first problem that continuum observers will notice with their EVLA data is that the spectral and time resolution of the data, by default anyway, will be rather more than their science requires. It will be possible to instruct the software which extracts data from the archive to do some averaging in both frequency and time. However, detailed editing for RFI and other issues may require excellent resolution in both these domains. After the data have been edited, you can average data in both domains so long as you are careful not to average so much that you produce radial (bandwidth) and/or transverse (time) smearing within the image area. Note that the increased sensitivity of the EVLA will increase the area over which non-negligible astronomical objects may be found while the wide bandwidth will mean that lowest frequency part of your band will be sensitive, because of its larger primary beam, to a much larger area on the sky than the highest frequency part. The spectral averaging can be done with SPLIT; use APARM(1)=1 and set NCHAV, CHINC, and perhaps SMOOTH appropriately. Similarly, AVSPC can be used with AVOPTION=’SUBS’, setting CHANNEL and SMOOTH suitably. You will almost certainly wish to retain some spectral separation, so do not use the “channel 0” option. Time averaging should be done with UVAVG: <source lang="text"> > DEFAULT UVAVG ; INP to reset all adverbs and choose the task. > INDI Sn; GETN Sm CR to select the calibrated target data set on disk Sn and catalog number Sm. > YINC ∆t CR to average to ∆t seconds. > GO CR to produce the averaged data set. </source> UBAVG will do a more aggressive averaging, using baseline-dependent time intervals appropriate for the desired field of view. Do not use UBAVG if you are planning to use self-calibration since it destroys the time regularity in the data on which CALIB depends. IMAGR may now do this extra averaging for you on the fly to reduce the size of the work file it uses. Set IM2PARM(11) and (12). Imaging of the continuum is discussed at great length in Chapter 5 and those details will not be repeated here. Bandwidth-synthesis imaging, which will be the only form of continuum imaging with the EVLA, will make certain adverbs more important. Set BCHAN and ECHAN to avoid the noisier edge channels. Set NCHAV = ECHAN - BCHAN + 1 and CHINC = NCHAV. This will then image all of your IFs and spectral channels into a single image, positioning each channel correctly in the uv plane. With the EVLA, you will be imaging a wider field of view than you did with the VLA. Use SETFC with IMSIZE 0 ; CELLSIZE 0 to see if you should image with a single facet or with multiple facets. If using multiple facets and trying for significant dynamic range, start imaging with OVERLAP 2 ; ONEBEAM -1, but consider OVRSWTCH = -0.05 or so to switch into faster methods of Cleaning when the dynamic range in the residual is small enough. 31DEC09 and later versions of IMAGR allow you to request automatic finding of the Clean boxes (IM2PARM of 1 through 6). In cases with low sidelobes, this works rather well, but you should probably keep an eye on what it does with DOTV 1 in any case. IM2PARM(12) controls the baseline-dependent time averaging while specifying the maximum field of view you expect. This allows you to reduce the size of the work file considerably which will at least reduce the time required for many of the steps in the imaging proportionally. It may be rather better than that if the work file is very large otherwise, requiring actual reading of the disk every time the data are accessed. Note, however, that the uniform weighting of your data will be affected. This averaging reduces the number of samples at short spacings disproportionally and so appears to reduce their weight in the imaging. Some UVTAPER could be reduce to compensate for this. By default, bandwidth synthesis imaging assumes that the primary beam and all continuum sources are the same at every frequency. In fact, the primary beam size varies linearly with frequency (to first order anyway) and sources have spectral index. IMAGR will allow you to compensate for the average spectral index at almost no cost with IMAGRPRM(2). A far more accurate and expensive correction for spectral index may be made if you do the following. First image each spectral channel (or group of closely-spaced channels) separately. Combine them into a cube with FQUBE, transpose the cube with TRANS, and solve for spectral index images with SPIXR. To use these images, set IMAGRPRM(17) to a radius (> 0) in pixels of a smoothing area and put the image name parameters in the 3rd and 4th input image names. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. The change of primary beam with frequency may be corrected by setting IMAGRPRM(1) = 25 for the diameter of the EVLA dishes. Note that this algorithm is expensive, but that it can be sped up with judicious use of the FQTOL parameter. These two corrections work together, so that doing both costs very little more than doing just one of them. If you are observing a strong source and trying for very high dynamic range, you will probably have to correct for errors that are baseline- rather than antenna-dependent. One source of these errors is the antenna polarization leakage which affects the parallel-hand visibilities in a non-closing fashion. Task BLCAL can be used after you have as good an image as you can get without it. This task will divide the data by the model and average over a user-specified time to find baseline-dependent corrections which may then be applied to the data by setting adverb BLVER. We recommend that you average the divided data over all of the times in your data to get a single correction for each baseline (and IF and polarization). If you use shorter intervals, you run the risk of forcing your data to look too much like your model. Since the polarization leakage is probably a function of frequency, an experimental version of BLCAL called BLCHN has been released. It determines the same correction but does not average over channels. The correction is saved in a table which POSSM and BPLOT are able to display. However, the calibration routines do not know how to apply this table, so BLCHN write out the corrected data as well as the table. == Concluding Remarks, References, Pictures == AIPS itself, and particularly this appendix, do not begin to cover all of the issues that will arise with EVLA data. The increased bandwidth will probably cause the polarization calibration to change from one complex number per antenna per IF to a complete complex spectrum much like a bandpass. This will force major revisions to the AIPS code which deals with this area. See § 4.6 for information about current AIPS methods of polarization calibration, which average all spectral channels within an IF except for those which are flagged. The increased sensitivity of the EVLA will means that imaging will no longer be able to ignore effects that are difficult to correct such as pointing errors, beam squint, variable antenna polarization across the field, leakage of polarized signal into the parallel-hand visibilities, etc., etc. These are research topics which may have solutions in other software packages such as OBIT and CASA. [[Image:OrionA_Kspectrum.png||The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten.]] Figure E.1: The spectrum of the hot core of Orion A at K band. Three separate observations of 8192 channels each 0.125 MHz wide were made using 12 antennas in the D array. Two hours total telescope time went into each of the two lower thirds of the spectrum and 1 hour was used for the highest third. The plot was made using ISPEC over a 54 by 60 arc second area. Line identifications provided by Karl Menten. 7ce0cf68d2b0209d4137ca14b48f79f5f129374d 6 6 17 2010-04-12T23:30:01Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 10 7 42 2010-04-13T15:29:14Z Jmcmulli 2 Created page with '{|style="background-color:#CCFFCC;" | [[EVLA Reduction Strategy|Special Considerations for EVLA Data Calibration and Imaging in AIPS]] | [[Image:OrionA_Kspectrum.png|thumb]] |}' wikitext text/x-wiki {|style="background-color:#CCFFCC;" | [[EVLA Reduction Strategy|Special Considerations for EVLA Data Calibration and Imaging in AIPS]] | [[Image:OrionA_Kspectrum.png|thumb]] |} 681653e90d2146e8d9769614e6d2edb4ff27576a 43 42 2010-04-13T15:31:17Z Jmcmulli 2 wikitext text/x-wiki {|style="background-color:#EAF5FB;" | [[EVLA Reduction Strategy|Special Considerations for EVLA Data Calibration and Imaging in AIPS]] | [[Image:OrionA_Kspectrum.png|thumb]] |} c54352c2d6c0486484be36d46f8b2f69aa665039 49 43 2010-04-13T17:16:05Z Jmcmulli 2 wikitext text/x-wiki {|style="background-color:#EAF5FB;" | [[:Category:Post-Processing|EVLA Reduction Strategy|Special Considerations for EVLA Data Calibration and Imaging in AIPS]] ]] | [[Image:OrionA_Kspectrum.png|thumb]] |} c665c477962277fafa05589242d9535a152222fe 50 49 2010-04-13T17:16:40Z Jmcmulli 2 wikitext text/x-wiki {|style="background-color:#EAF5FB;" | [[:Category:Post-Processing|EVLA Reduction Strategy|Special Considerations for EVLA Data Calibration and Imaging in AIPS ]] | [[Image:OrionA_Kspectrum.png|thumb]] |} 6f37b7d11afeee70ca204f10ea73147d5d3aa345 10 7 51 50 2010-04-13T17:17:52Z Jmcmulli 2 wikitext text/x-wiki {|style="background-color:#EAF5FB;" | [[:Category:Post-Processing | EVLA Reduction Strategy | Special Considerations for EVLA Data Calibration and Imaging in AIPS ]] | [[Image:OrionA_Kspectrum.png|thumb]] |} 6f874b1018b15075555514f5fcafbe08816f1dd1 10 2 52 46 2010-04-13T17:19:05Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] |} 773f85d4a8c3e1395486552371d9c4944e3ef3c4 14 8 53 2010-04-13T18:01:03Z Jmcmulli 2 Created page with '= Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference…' wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. ec5f2b7fdea3734d42bf89e9d8e1f22ba142c3be 54 53 2010-04-13T18:10:45Z Jmcmulli 2 wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>Delta I</math> where: SEF D ∆Im= 􏰁 (1) ηc npolN(N − 1)tint∆ν – SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the VLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. – ηc is the correlator efficiency (at least 0.92 for the EVLA). – npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. – N is the number of antennas. – tint is the total on-source integration time in seconds. – ∆ν is the bandwidth in Hz. e545d6a9db2924c113054d3c629d45f2cfdbeff3 55 54 2010-04-13T19:35:04Z Jmcmulli 2 /* Sensitivity */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> feea06f4bfc4c50759657502ff419471a77fd5de 56 55 2010-04-13T20:47:18Z Jmcmulli 2 /* Expected Capabilities: Receivers */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png|left]] Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> 1661971ef98ee817beec8721348fb7c6bef2125c 58 56 2010-04-13T20:49:26Z Jmcmulli 2 /* Expected Capabilities: Receivers */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> 1b978ebdcd756b944e114b687a790c061dcbfaf3 59 58 2010-04-13T22:08:31Z Jmcmulli 2 /* Sensitivity */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: b93a2d808c8e7177a7f2c75bbafb69d66ce6477b 60 59 2010-04-13T22:09:46Z Jmcmulli 2 wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. 10d881113640ac6492d4e65430c8780e591b1fad 61 60 2010-04-13T22:20:33Z Jmcmulli 2 wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == == Data Processing == == Travel Support for Visiting the DSOC and EVLA == == Student Assistance for Data Reduction Visits to the DSOC == == Computing at the DSOC == == Reservations for the EVLA site and/or DSOC == == Staying in Socorro == == Help for Visitors to the EVLA and DSOC == == On-Line Information about the NRAO and the EVLA == = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === === Dissertations === === Preprints === === Reprints === === Page Charge Support === = Documentation = = Key Personnel = = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. adbde5e2208f53f7a06b03221a9cf526f83bea7e 62 61 2010-04-13T22:37:35Z Jmcmulli 2 /* Key Personnel */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == == Data Processing == == Travel Support for Visiting the DSOC and EVLA == == Student Assistance for Data Reduction Visits to the DSOC == == Computing at the DSOC == == Reservations for the EVLA site and/or DSOC == == Staying in Socorro == == Help for Visitors to the EVLA and DSOC == == On-Line Information about the NRAO and the EVLA == = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === === Dissertations === === Preprints === === Reprints === === Page Charge Support === = Documentation = = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 7a090400160738958820a05f2c514b3bd76c77d2 63 62 2010-04-13T22:40:59Z Jmcmulli 2 /* Documentation */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == == Data Processing == == Travel Support for Visiting the DSOC and EVLA == == Student Assistance for Data Reduction Visits to the DSOC == == Computing at the DSOC == == Reservations for the EVLA site and/or DSOC == == Staying in Socorro == == Help for Visitors to the EVLA and DSOC == == On-Line Information about the NRAO and the EVLA == = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === === Dissertations === === Preprints === === Reprints === === Page Charge Support === = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 2396a8f1d7a054fd19d6772dafd27f364361e2b7 64 63 2010-04-13T22:42:39Z Jmcmulli 2 /* Page Charge Support */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == == Data Processing == == Travel Support for Visiting the DSOC and EVLA == == Student Assistance for Data Reduction Visits to the DSOC == == Computing at the DSOC == == Reservations for the EVLA site and/or DSOC == == Staying in Socorro == == Help for Visitors to the EVLA and DSOC == == On-Line Information about the NRAO and the EVLA == = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === === Dissertations === === Preprints === === Reprints === === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 7d37306eef6f63862e5065b8ab67fa8c2c0d5517 65 64 2010-04-13T22:43:00Z Jmcmulli 2 /* Reprints */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == == Data Processing == == Travel Support for Visiting the DSOC and EVLA == == Student Assistance for Data Reduction Visits to the DSOC == == Computing at the DSOC == == Reservations for the EVLA site and/or DSOC == == Staying in Socorro == == Help for Visitors to the EVLA and DSOC == == On-Line Information about the NRAO and the EVLA == = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === === Dissertations === === Preprints === === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. ba95d1bc37e5af04fd6e91355d24d764fcb57ace 66 65 2010-04-13T22:43:24Z Jmcmulli 2 /* Preprints */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == == Data Processing == == Travel Support for Visiting the DSOC and EVLA == == Student Assistance for Data Reduction Visits to the DSOC == == Computing at the DSOC == == Reservations for the EVLA site and/or DSOC == == Staying in Socorro == == Help for Visitors to the EVLA and DSOC == == On-Line Information about the NRAO and the EVLA == = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === === Dissertations === === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. d80e4a762c7d7aeb3e50644e1086da1bd4b0017f 67 66 2010-04-13T22:43:52Z Jmcmulli 2 /* Dissertations */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == == Data Processing == == Travel Support for Visiting the DSOC and EVLA == == Student Assistance for Data Reduction Visits to the DSOC == == Computing at the DSOC == == Reservations for the EVLA site and/or DSOC == == Staying in Socorro == == Help for Visitors to the EVLA and DSOC == == On-Line Information about the NRAO and the EVLA == = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. d5c46ac37d7da52d91333c09895f07f0dce55ee1 68 67 2010-04-13T22:44:36Z Jmcmulli 2 /* Acknowledgement to NRAO */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == == Data Processing == == Travel Support for Visiting the DSOC and EVLA == == Student Assistance for Data Reduction Visits to the DSOC == == Computing at the DSOC == == Reservations for the EVLA site and/or DSOC == == Staying in Socorro == == Help for Visitors to the EVLA and DSOC == == On-Line Information about the NRAO and the EVLA == = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 2686ded3bf9034fb108799e1af9ee261420ba29a 69 68 2010-04-13T22:45:06Z Jmcmulli 2 /* On-Line Information about the NRAO and the EVLA */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == == Data Processing == == Travel Support for Visiting the DSOC and EVLA == == Student Assistance for Data Reduction Visits to the DSOC == == Computing at the DSOC == == Reservations for the EVLA site and/or DSOC == == Staying in Socorro == == Help for Visitors to the EVLA and DSOC == == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 8d632d5de7a718bb2243faa58a67d34301645b76 70 69 2010-04-13T22:45:42Z Jmcmulli 2 /* Help for Visitors to the EVLA and DSOC */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == == Data Processing == == Travel Support for Visiting the DSOC and EVLA == == Student Assistance for Data Reduction Visits to the DSOC == == Computing at the DSOC == == Reservations for the EVLA site and/or DSOC == == Staying in Socorro == == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. e8b3cd143f1f5798ff9bf65fd8f367134185f408 71 70 2010-04-13T22:46:14Z Jmcmulli 2 /* Staying in Socorro */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == == Data Processing == == Travel Support for Visiting the DSOC and EVLA == == Student Assistance for Data Reduction Visits to the DSOC == == Computing at the DSOC == == Reservations for the EVLA site and/or DSOC == == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 572d0db6500d2f2fa3f25f1e5376e6dc29afad7c 72 71 2010-04-13T22:47:01Z Jmcmulli 2 /* Reservations for the EVLA site and/or DSOC */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == == Data Processing == == Travel Support for Visiting the DSOC and EVLA == == Student Assistance for Data Reduction Visits to the DSOC == == Computing at the DSOC == == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. ee9f73d5b46fd9e072902207a952776d636b113c 73 72 2010-04-13T22:47:40Z Jmcmulli 2 /* Computing at the DSOC */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == == Data Processing == == Travel Support for Visiting the DSOC and EVLA == == Student Assistance for Data Reduction Visits to the DSOC == == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 573edf73f20a2438163ebc72be72e6a3364b6474 74 73 2010-04-13T22:48:03Z Jmcmulli 2 /* Student Assistance for Data Reduction Visits to the DSOC */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == == Data Processing == == Travel Support for Visiting the DSOC and EVLA == == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 0600e33cbb8d49ec42d664c69d203679eb48c2c3 75 74 2010-04-13T22:49:29Z Jmcmulli 2 /* Travel Support for Visiting the DSOC and EVLA */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == == Data Processing == == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 7cde53c88ad695fdab82fd555e7a367500919c2b 76 75 2010-04-13T22:50:02Z Jmcmulli 2 /* Data Processing */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 6dabe8e37ba4bc84121acdae65056a3dbf877cca 77 76 2010-04-13T22:51:59Z Jmcmulli 2 /* Data Access */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. ff6b40df236089666548cd7dca54d684c85c6651 78 77 2010-04-13T22:52:20Z Jmcmulli 2 /* The Observations and Remote Observing */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 150b7306a7641d9057d6f16364d5c07efa9a1846 79 78 2010-04-13T22:52:50Z Jmcmulli 2 /* Fixed Date and Dynamic Scheduling */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 85f9db254c7e8a4f86144837683f333567863f02 80 79 2010-04-13T22:54:26Z Jmcmulli 2 /* Observation Preparation */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. b0c8ebba0f3fb75593a9767b001991de1ffa536e 81 80 2010-04-13T22:54:55Z Jmcmulli 2 /* Data Analysts and General Assistance */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 366c6cefa1bb03d21ce8c268b29102364428c242 82 81 2010-04-13T22:56:10Z Jmcmulli 2 /* Rapid Response Science */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 1a250601ac8bb928775c4af66196211f399d8859 83 82 2010-04-13T22:57:12Z Jmcmulli 2 /* Obtaining Observing Time on the EVLA */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 7af0054c4ff3b8711fe6e8dfb01becec1f30dd08 84 83 2010-04-13T22:57:38Z Jmcmulli 2 /* Pulsar Observing */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. e30a68abe7e2eb0bdc6cff6e09bbefe065d3e41b 85 84 2010-04-13T22:58:11Z Jmcmulli 2 /* Combining Configurations and Mosaicing */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. e3f865739208de50dbcd62ac7b58bc7d48052790 86 85 2010-04-13T22:59:21Z Jmcmulli 2 /* Shadowing and Cross-Talk */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 9363217949c7ebc43e6fc36a7a74e48aebb6fd2e 87 86 2010-04-13T23:00:06Z Jmcmulli 2 /* Snapshots */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 1fb1f2c664d8b4e71b2087916a168d8d1406d399 88 87 2010-04-13T23:00:26Z Jmcmulli 2 /* VLBI Observations */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 4434e5e545714b8f4370f00a667d4997502eb9a8 89 88 2010-04-13T23:02:06Z Jmcmulli 2 /* Correlator Configurations */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === == Polarization == == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 8cc780ad67568280ed4de11df1a8e8159a15429a 90 89 2010-04-14T14:04:58Z Jmcmulli 2 /* Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 68b71dede78076876e84035e5268a1eeb11e4b0b 91 90 2010-04-14T14:05:58Z Jmcmulli 2 /* Polarization */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 355504fdf30f55eecf464b3ef1bfd6b8fc44d3aa 92 91 2010-04-14T14:22:05Z Jmcmulli 2 /* Correlator Configurations */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 463e1400154df849b8a4abf88fcf911bb0bf2506 93 92 2010-04-14T14:48:46Z Jmcmulli 2 /* General Guidelines for Gain Calibration */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 1145c6951e463325b626df258559d15c67a9fde9 94 93 2010-04-14T14:54:43Z Jmcmulli 2 /* Calibration and the Flux Density Scale */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. c2de485a53eaa8a0c51643fc220813657bd5dcfb 95 94 2010-04-14T14:55:06Z Jmcmulli 2 /* Sidelobes from Strong Sources */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 28bf9106c166a6779edc2b20945201dc01c8d5fb 96 95 2010-04-14T14:55:28Z Jmcmulli 2 /* Sidelobes from Confusing Sources */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. a4041fcb092f0b6d1a3112e429137e73d4d88afe 97 96 2010-04-14T14:55:58Z Jmcmulli 2 /* Poorly Sampled Fourier Plane */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 8f7474e69857746ca589f7253907f8bf11755139 98 97 2010-04-14T14:56:16Z Jmcmulli 2 /* Invisible Structures */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == === Image Fidelity === === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 4a14008d0275ffc2d4a8db8e7daaaab0b39b11a3 99 98 2010-04-14T15:00:59Z Jmcmulli 2 /* Limitations on Imaging Performance */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 012903ae9457ea422b23a07570b5edd7a4ff3371 6 9 57 2010-04-13T20:47:46Z Jmcmulli 2 Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency c wikitext text/x-wiki Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. b5123426ceadb1e9463aac895961d637988a2209 6 10 100 2010-04-14T15:01:24Z Jmcmulli 2 This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. wikitext text/x-wiki This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. c970dfb1b339f5e2c1c961b835ed714e275b3839 14 8 101 99 2010-04-14T15:02:41Z Jmcmulli 2 /* Image Fidelity */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 3c1daef0c13ef9bd53cf0cbe6f612c2f4b338898 102 101 2010-04-14T15:05:16Z Jmcmulli 2 /* Positional Accuracy */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 45a6c4a0cca52aa149f826fba0100722b0d12cfd 104 102 2010-04-14T15:09:13Z Jmcmulli 2 /* Subarrays */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. bd4d8943675b5071e959dbc00b882fd767210220 105 104 2010-04-14T15:17:18Z Jmcmulli 2 /* Radio-Frequency Interference */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 553d8b133211ac0cffd483322405494a467e5540 106 105 2010-04-14T15:19:19Z Jmcmulli 2 /* Time Resolution and Data Rates */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. e6b12eb8f4c57dff6eb43b83da62d868ffde0bdf 107 106 2010-04-14T15:19:50Z Jmcmulli 2 /* Non-Coplanar Baselines */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. ef1bcdd70affd8d713601e1d4a32ad93bdd71445 108 107 2010-04-14T15:25:33Z Jmcmulli 2 /* Time-Averaging Loss */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="2"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. d049e491a8b7e56315057d68bd7bea008841ae01 109 108 2010-04-14T15:25:58Z Jmcmulli 2 /* Time-Averaging Loss */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 6a9888ddb237050a56a92ff97052adb38c3bb0b8 110 109 2010-04-14T15:29:16Z Jmcmulli 2 /* Chromatic Aberration (Bandwidth Smearing) */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochro- matic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :'''Note:''' The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ0/θHPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. f414d5251e523b37332b83f42f1051aa0b401689 111 110 2010-04-14T15:30:15Z Jmcmulli 2 /* Primary Beam */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the indi- vidual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochro- matic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :'''Note:''' The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ0/θHPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. fd57762a69ecfc904ebb8d4f2eec4863b7063e20 112 111 2010-04-14T15:30:42Z Jmcmulli 2 /* Field of View */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the indi- vidual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochro- matic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :'''Note:''' The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ0/θHPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. df96b05908eb0217770ef0bc48ec937bcb0c7a2e 113 112 2010-04-14T15:36:14Z Jmcmulli 2 /* EVLA Frequency Bands and Tunability */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" |+ '''Table 7: Default frequencies for "continuum" applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :'''Notes:''' :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference- free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1860 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the indi- vidual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochro- matic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :'''Note:''' The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ0/θHPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 505dac0f1e00b27892540420f07839f552401dba 114 113 2010-04-14T15:40:39Z Jmcmulli 2 /* Sensitivity */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\alpha^2+\beta^2=1</math> where: :– SEFD is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEF D indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62Tsys/ηA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band. :– ηc is the correlator efficiency (at least 0.92 for the EVLA). :– npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– tint is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the SEFD as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the SEFD is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard C configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in D configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the D configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: [[Image:BrightnessTemp.png]] where Tb is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in Section 7. For observers interested in Hi in galaxies, a number of interest is the sensitivity of the observation to the Hi mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): [[Image:MHI.png]] where D is the distance to the galaxy in M_pc, and S∆V is the Hi line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" |+ '''Table 7: Default frequencies for "continuum" applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :'''Notes:''' :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference- free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1860 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the indi- vidual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochro- matic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :'''Note:''' The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ0/θHPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. a5c00114c0731b8c60d166ad03ab5256c4738c8a 117 114 2010-04-14T15:49:53Z Jmcmulli 2 /* Sensitivity */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: [[Image:thermalnoise.png]] where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. [[Image:SEFD.png]] {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: [[Image:BrightnessTemp.png]] where Tb is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in Section 7. For observers interested in Hi in galaxies, a number of interest is the sensitivity of the observation to the Hi mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): [[Image:MHI.png]] where D is the distance to the galaxy in M_pc, and S∆V is the Hi line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" |+ '''Table 7: Default frequencies for "continuum" applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :'''Notes:''' :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference- free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1860 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the indi- vidual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochro- matic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :'''Note:''' The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ0/θHPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 4e3b91ef220decf4f9d4474d8bd6e82a9136ecc0 120 117 2010-04-14T15:56:53Z Jmcmulli 2 /* Sensitivity */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: [[Image:thermalnoise.png]] where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. [[Image:SEFD.png|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] : '''Figure 2:''' ''SEFD'' for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: [[Image:BrightnessTemp.png]] where Tb is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in Section 7. For observers interested in Hi in galaxies, a number of interest is the sensitivity of the observation to the Hi mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): [[Image:MHI.png]] where D is the distance to the galaxy in M_pc, and S∆V is the Hi line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" |+ '''Table 7: Default frequencies for "continuum" applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :'''Notes:''' :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference- free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1860 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the indi- vidual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochro- matic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :'''Note:''' The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ0/θHPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 0f7190233ea42c12c1a1268ced3c85588e0811ac 121 120 2010-04-14T16:00:13Z Jmcmulli 2 /* EVLA Early Science */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || D || C || B |- | 2011 || A || D || C |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] :'''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA con- struction project began. The OSRO program extends this into the EVLA era by providing early access to a number of correlator capabilities and observing modes for the general user community after the VLA correlator has been turned off in January 2010. These modes represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in Section 4. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} :'''Note:''' The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in So- corro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific pro- ductivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astro- nomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: [[Image:thermalnoise.png]] where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. [[Image:SEFD.png|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] : '''Figure 2:''' ''SEFD'' for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: [[Image:BrightnessTemp.png]] where Tb is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in Section 7. For observers interested in Hi in galaxies, a number of interest is the sensitivity of the observation to the Hi mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): [[Image:MHI.png]] where D is the distance to the galaxy in M_pc, and S∆V is the Hi line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" |+ '''Table 7: Default frequencies for "continuum" applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :'''Notes:''' :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference- free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1860 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the indi- vidual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochro- matic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :'''Note:''' The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ0/θHPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 261fb32551dd096cf32360ca69f4c82bc8f25e31 122 121 2010-04-14T16:01:44Z Jmcmulli 2 /* An Overview of the EVLA */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' || '''B''' |- | 2011 || '''A''' || '''D''' || '''C''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] :'''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA con- struction project began. The OSRO program extends this into the EVLA era by providing early access to a number of correlator capabilities and observing modes for the general user community after the VLA correlator has been turned off in January 2010. These modes represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in Section 4. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. {| border="1" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} :'''Note:''' The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in So- corro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific pro- ductivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astro- nomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: [[Image:thermalnoise.png]] where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. [[Image:SEFD.png|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] : '''Figure 2:''' ''SEFD'' for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear. {| border="1" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: [[Image:BrightnessTemp.png]] where Tb is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in Section 7. For observers interested in Hi in galaxies, a number of interest is the sensitivity of the observation to the Hi mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): [[Image:MHI.png]] where D is the distance to the galaxy in M_pc, and S∆V is the Hi line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" |+ '''Table 7: Default frequencies for "continuum" applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :'''Notes:''' :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference- free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1860 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the indi- vidual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochro- matic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :'''Note:''' The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ0/θHPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 6c16f06020b14c056ff10cd4ddd7b3e5f49506d4 123 122 2010-04-14T16:05:32Z Jmcmulli 2 wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' || '''B''' |- | 2011 || '''A''' || '''D''' || '''C''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] :'''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA con- struction project began. The OSRO program extends this into the EVLA era by providing early access to a number of correlator capabilities and observing modes for the general user community after the VLA correlator has been turned off in January 2010. These modes represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in Section 4. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} :'''Note:''' The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in So- corro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific pro- ductivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astro- nomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: [[Image:thermalnoise.png]] where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. [[Image:SEFD.png|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] : '''Figure 2:''' ''SEFD'' for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear. {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (\mu Jy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: [[Image:BrightnessTemp.png]] where Tb is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in Section 7. For observers interested in Hi in galaxies, a number of interest is the sensitivity of the observation to the Hi mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): [[Image:MHI.png]] where D is the distance to the galaxy in M_pc, and S∆V is the Hi line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for "continuum" applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :'''Notes:''' :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference- free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1860 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the indi- vidual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochro- matic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :'''Note:''' The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ0/θHPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. fbffcc6a1047e9a0e29bdf09a268608c5d7ee279 124 123 2010-04-14T16:08:02Z Jmcmulli 2 /* Sensitivity */ wikitext text/x-wiki = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' || '''B''' |- | 2011 || '''A''' || '''D''' || '''C''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] :'''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA con- struction project began. The OSRO program extends this into the EVLA era by providing early access to a number of correlator capabilities and observing modes for the general user community after the VLA correlator has been turned off in January 2010. These modes represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in Section 4. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} :'''Note:''' The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in So- corro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific pro- ductivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astro- nomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: [[Image:thermalnoise.png]] where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. [[Image:SEFD.png|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] : '''Figure 2:''' ''SEFD'' for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear. {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 600 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: [[Image:BrightnessTemp.png]] where Tb is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in Section 7. For observers interested in Hi in galaxies, a number of interest is the sensitivity of the observation to the Hi mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): [[Image:MHI.png]] where D is the distance to the galaxy in M_pc, and S∆V is the Hi line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for "continuum" applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :'''Notes:''' :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference- free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1860 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the indi- vidual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochro- matic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :'''Note:''' The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ0/θHPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 468c933780b54d172dfa17fe6f3cab223368069a 126 124 2010-04-20T15:51:26Z Cchandle 6 Added title and version date; updated L-band default frequencies; fixed Ka-band SEFD in table 6. wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: April 20, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. == VLA to EVLA Transition == The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Full EVLA correlator installation || 2010 Q2 |- | Last antenna retrofitted || 2010 Q2 |- | Last receiver installed || 2012 Q3 |- |} = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside- down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are de- ployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Feb-May ! Jun-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' || '''B''' |- | 2011 || '''A''' || '''D''' || '''C''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See Section 4.15 for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands. The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities: Antennas == Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010. == Expected Capabilities: Receivers == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] :'''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA con- struction project began. The OSRO program extends this into the EVLA era by providing early access to a number of correlator capabilities and observing modes for the general user community after the VLA correlator has been turned off in January 2010. These modes represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in Section 4. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} :'''Note:''' The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in So- corro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific pro- ductivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astro- nomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half- power), and the scale at which severe attenuation of large scale structure occurs. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: [[Image:thermalnoise.png]] where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth theta_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale theta_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, theta_HPBW is the synthesized beam width (FWHM), and theta_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. [[Image:SEFD.png|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] : '''Figure 2:''' ''SEFD'' for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear. {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: [[Image:BrightnessTemp.png]] where Tb is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in Section 7. For observers interested in Hi in galaxies, a number of interest is the sensitivity of the observation to the Hi mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): [[Image:MHI.png]] where D is the distance to the galaxy in M_pc, and S∆V is the Hi line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for "continuum" applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :'''Notes:''' :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference- free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the indi- vidual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochro- matic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :'''Note:''' The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ0/θHPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. cd32620809e4d268cce68c99a06045530f1ffdff 136 126 2010-05-11T21:10:41Z Cchandle 6 Updates for the 2010 May 15 proposal call. wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: April 20, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[Key Personnel]], or refer to the manuals and documentation listed in [[Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q3 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Jul ! Aug-Nov ! Dec-Feb |- | 2010 || '''D''' || '''C''' || '''B''' |- | 2011 || '''A''' || '''D''' || '''C''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim we plan to mount and test the compatibility of the existing 74-MHz dipoles with the wideband EVLA electronics in the upcoming C-configuration, with the goal of providing a low frequency observing capability in the B/BnA/A configurations. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] :'''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} :'''Note:''' The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[Documentation]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> [[Image:thermalnoise.png]] where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[Image:SEFD.png|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] : '''Figure 2:''' ''SEFD'' for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear. {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: [[Image:BrightnessTemp.png]] where Tb is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in Section 7. For observers interested in Hi in galaxies, a number of interest is the sensitivity of the observation to the Hi mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): [[Image:MHI.png]] where D is the distance to the galaxy in M_pc, and S∆V is the Hi line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for "continuum" applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :'''Notes:''' :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference- free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the indi- vidual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochro- matic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :'''Note:''' The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ0/θHPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} == Subarrays == The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large- scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow- bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars. Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63 |- | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62 |- | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11 |- | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58 |- | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41 |- | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22 |- | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22 |- |} For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geosta- tionary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instru- mental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 3ad06efff36dea0ae3fc5954f438a3c0916088a9 137 136 2010-05-11T22:59:11Z Cchandle 6 wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 11, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q3 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Jul ! Aug-Nov ! Dec-Feb |- | 2010 || '''D''' || '''C''' || '''B''' |- | 2011 || '''A''' || '''D''' || '''C''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim we plan to mount and test the compatibility of the existing 74-MHz dipoles with the wideband EVLA electronics in the upcoming C-configuration, with the goal of providing a low frequency observing capability in the B/BnA/A configurations. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[image:EVLA_WidebandRxAvailability.png]] :'''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown. Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, June 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5 |- | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27 |- |} :'''Note:''' The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. Footnotes: :1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. :2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. :4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. :5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). :6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[Image:SEFD.png|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] : '''Figure 2:''' ''SEFD'' for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear. {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :'''Notes:''' :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference- free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :'''Note:''' The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660 – 1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68– 10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]] :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB). [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]] :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB). {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3). :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Data Analysts and General Assistance == General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Miscellaneous = == Publication Guidelines == === Acknowledgement to NRAO === Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.'' === Dissertations === Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. === Preprints === NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). === Reprints === Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. === Page Charge Support === The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see Section 6.1.1). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 4f20d57ca421e2a66edebbc9944bd64ac2fb81f1 139 137 2010-05-12T07:50:16Z Cchandle 6 More updates for the May 2010 call for proposals. wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 11, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q3 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Jul ! Aug-Nov ! Dec-Feb |- | 2010 || '''D''' || '''C''' || '''B''' |- | 2011 || '''A''' || '''D''' || '''C''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim we plan to mount and test the compatibility of the existing 74-MHz dipoles with the wideband EVLA electronics in the upcoming C-configuration, with the goal of providing a low frequency observing capability in the B/BnA/A configurations. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, July 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 8 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 23 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 1 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 4 || align='center'| 4 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 24 || align='center'| 24 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 4d87eb61af8b5c47aaeac859ecae20f3b473bfab 140 139 2010-05-12T14:46:16Z Cchandle 6 wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 12, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q3 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Jul ! Aug-Nov ! Dec-Feb |- | 2010 || '''D''' || '''C''' || '''B''' |- | 2011 || '''A''' || '''D''' || '''C''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim we plan to mount and test the compatibility of the existing 74-MHz dipoles with the wideband EVLA electronics in the upcoming C-configuration, with the goal of providing a low frequency observing capability in the B/BnA/A configurations. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, July 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 8 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 23 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 1 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 4 || align='center'| 4 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 24 || align='center'| 24 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|none|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. e12b0c2e01ae598842fd41aeda3be37225b0d74a 6 11 103 2010-04-14T15:08:02Z Jmcmulli 2 Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, unca wikitext text/x-wiki Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. 32f83db50900f0333550e2e7ea248971b8f61a82 6 12 115 2010-04-14T15:40:57Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 13 116 2010-04-14T15:41:35Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 14 118 2010-04-14T15:50:09Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 15 119 2010-04-14T15:50:52Z Jmcmulli 2 Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, a wikitext text/x-wiki Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear. 73d0202b11dfa76074665912604c2f3fb1f47a52 10 2 125 52 2010-04-14T16:09:41Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] |} 8ee9e58caec8bd424c1ac5ff250e6e3c23488667 127 125 2010-04-21T22:33:44Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] |} b7954e8a8c42b652bd1e28110e5009987d5ec722 141 127 2010-06-02T19:03:39Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] · [[EVLA Antenna-Receiver Availability Table]] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] |} 88bd3bda8b6ed8e5c5ce5e182f7bf6ca7ab9546f 144 141 2010-06-28T00:35:21Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] · [https://staff.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [http://mctest.evla.nrao.edu/cgi-bin/evla/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] |} 4e56191cb752f767e3763c9d6e260dee7097858f 147 144 2010-06-29T19:19:46Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [http://mctest.evla.nrao.edu/cgi-bin/evla/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] |} 2598b356cb398fba28dc21a7a7bd254ab9f350a2 148 147 2010-07-19T20:52:39Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [http://mctest.evla.nrao.edu/cgi-bin/evla/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] · [[Key to Calcodes]] |} cfa457310ee81890daabb556c54517e23e96096b 0 16 128 2010-04-21T22:40:49Z Jmcmulli 2 Created page with '= A = {| border="1" |+ A ! Acronym ! Meaning |- |} = B = {| border="1" |+ B ! Acronym ! Meaning |- |} {| border="1" |+ C ! Acronym ! Meaning |- |} {| |+ D ! Acronym ! Meaning …' wikitext text/x-wiki = A = {| border="1" |+ A ! Acronym ! Meaning |- |} = B = {| border="1" |+ B ! Acronym ! Meaning |- |} {| border="1" |+ C ! Acronym ! Meaning |- |} {| |+ D ! Acronym ! Meaning |- |} {| |+ E ! Acronym ! Meaning |- |} {| |+ F ! Acronym ! Meaning |- |} {| |+ G ! Acronym ! Meaning |- |} {| |+ H ! Acronym ! Meaning |- |} {| |+ I ! Acronym ! Meaning |- |} {| |+ J ! Acronym ! Meaning |- |} {| |+ K ! Acronym ! Meaning |- |} {| |+ L ! Acronym ! Meaning |- |} {| |+ M ! Acronym ! Meaning |- |} {| |+ N ! Acronym ! Meaning |- |} {| |+ O ! Acronym ! Meaning |- |} {| |+ P ! Acronym ! Meaning |- |} {| |+ Q ! Acronym ! Meaning |- |} {| |+ R ! Acronym ! Meaning |- |} {| |+ S ! Acronym ! Meaning |- |} {| |+ T ! Acronym ! Meaning |- |} {| |+ U ! Acronym ! Meaning |- |} {| |+ V ! Acronym ! Meaning |- |} {| |+ W ! Acronym ! Meaning |- |} {| |+ X ! Acronym ! Meaning |- |} {| |+ Y ! Acronym ! Meaning |- |} {| |+ Z ! Acronym ! Meaning |- |} 856a66480fb95c5ff20d572f0020bc8a592b05a8 129 128 2010-04-21T22:56:30Z Jmcmulli 2 wikitext text/x-wiki = A B C = {| border="1" |+ A ! Acronym ! Meaning |- |} {| border="1" |+ B ! Acronym ! Meaning |- |} {| border="1" |+ C ! Acronym ! Meaning |- |} = D E F = {| border="1" |+ D ! Acronym ! Meaning |- | DTS || Digital Transmission System |- |} {| border="1" |+ E ! Acronym ! Meaning |- |} {| border="1" |+ F ! Acronym ! Meaning |- |} = G H I = {| border="1" |+ G ! Acronym ! Meaning |- |} {| border="1" |+ H ! Acronym ! Meaning |- |} {| border="1" |+ I ! Acronym ! Meaning |- |} = J K L = {| border="1" |+ J ! Acronym ! Meaning |- |} {| border="1" |+ K ! Acronym ! Meaning |- |} {| border="1" |+ L ! Acronym ! Meaning |- | LCP || Left Circular Polarization |- | LNA || Low Noise Amplifier |- | LO || Local Oscillator |- |} = M N O = {| border="1" |+ M ! Acronym ! Meaning |- | MIB || Module Interface Board |- |} {| border="1" |+ N ! Acronym ! Meaning |- |} {| border="1" |+ O ! Acronym ! Meaning |- |} = P Q R = {| border="1" |+ P ! Acronym ! Meaning |- |} {| border="1" |+ Q ! Acronym ! Meaning |- |} {| border="1" |+ R ! Acronym ! Meaning |- |} = S T U = {| border="1" |+ S ! Acronym ! Meaning |- |} {| border="1" |+ T ! Acronym ! Meaning |- |} {| border="1" |+ U ! Acronym ! Meaning |- |} = V W X = {| border="1" |+ V ! Acronym ! Meaning |- |} {| border="1" |+ W ! Acronym ! Meaning |- |} {| border="1" |+ X ! Acronym ! Meaning |- |} = Y Z = {| border="1" |+ Y ! Acronym ! Meaning |- |} {| border="1" |+ Z ! Acronym ! Meaning |- |} 499b8132980f23271f0d993a0c70bbb1fd9d0915 130 129 2010-04-23T21:59:38Z Jmcmulli 2 wikitext text/x-wiki = A B C = {| border="1" |+ A ! Acronym ! Meaning |- | ACU || Antenna Control Unit |- | AOC || Array Operations Center (now DSOC) |- |} {| border="1" |+ B ! Acronym ! Meaning |- |} {| border="1" |+ C ! Acronym ! Meaning |- | CDL || Central Development Laboratory |- | CMP || Control and Monitor Processor |- |} = D E F = {| border="1" |+ D ! Acronym ! Meaning |- | DCAF || Data Capture and Format |- | DSOC || Domenici Science Operations Center |- | DTS || Digital Transmission System |- | FRM || Focus Rotation Mount |} {| border="1" |+ E ! Acronym ! Meaning |- |} {| border="1" |+ F ! Acronym ! Meaning |- | FET || Field Effect Transistor |- |} = G H I = {| border="1" |+ G ! Acronym ! Meaning |- | GBT || Green Bank Telescope |- |} {| border="1" |+ H ! Acronym ! Meaning |- |} {| border="1" |+ I ! Acronym ! Meaning |- | IDCAF || interim Data Capture and Format |- | IPT || Integrated Product Team |- |} = J K L = {| border="1" |+ J ! Acronym ! Meaning |- |} {| border="1" |+ K ! Acronym ! Meaning |- |} {| border="1" |+ L ! Acronym ! Meaning |- | LCP || Left Circular Polarization |- | LNA || Low Noise Amplifier |- | LO || Local Oscillator |- |} = M N O = {| border="1" |+ M ! Acronym ! Meaning |- | MIB || Module Interface Board |- |} {| border="1" |+ N ! Acronym ! Meaning |- |} {| border="1" |+ O ! Acronym ! Meaning |- | OPT || Observation Preparation Tool |- | OTS || On-the sky |- | OMT || Ortho mode Transducer |- |} = P Q R = {| border="1" |+ P ! Acronym ! Meaning |- |} {| border="1" |+ Q ! Acronym ! Meaning |- |} {| border="1" |+ R ! Acronym ! Meaning |- | RFI || Radio Frequency Interference |- | RTCAT || Real-time Calibrator Analysis Tool |- |} = S T U = {| border="1" |+ S ! Acronym ! Meaning |- | SBC || Single Board Computer |- | SLC || Serial Line Controller |- |} {| border="1" |+ T ! Acronym ! Meaning |- |} {| border="1" |+ U ! Acronym ! Meaning |- |} = V W X = {| border="1" |+ V ! Acronym ! Meaning |- | VLA || Very Large Array |- | VLBA || Very Long Baseline Array |- | VME || |- | VPN || Virtual Private Network |- |} {| border="1" |+ W ! Acronym ! Meaning |- |} {| border="1" |+ X ! Acronym ! Meaning |- |} = Y Z = {| border="1" |+ Y ! Acronym ! Meaning |- |} {| border="1" |+ Z ! Acronym ! Meaning |- |} 25f7bdc7673a6762b667ba4748a71c399d89b31c 131 130 2010-04-24T03:07:46Z Jmcmulli 2 wikitext text/x-wiki = A B C = {| border="1" |+ A ! Acronym ! Meaning |- | ACU || Antenna Control Unit |- | AOC || Array Operations Center (now DSOC) |- |} {| border="1" |+ B ! Acronym ! Meaning |- |} {| border="1" |+ C ! Acronym ! Meaning |- | CDL || Central Development Laboratory |- | CMP || Control and Monitor Processor |- |} = D E F = {| border="1" |+ D ! Acronym ! Meaning |- | DCAF || Data Capture and Format |- | DSOC || Domenici Science Operations Center |- | DTS || Digital Transmission System |- | FRM || Focus Rotation Mount |} {| border="1" |+ E ! Acronym ! Meaning |- |} {| border="1" |+ F ! Acronym ! Meaning |- | FET || Field Effect Transistor |- |} = G H I = {| border="1" |+ G ! Acronym ! Meaning |- | GBT || Green Bank Telescope |- |} {| border="1" |+ H ! Acronym ! Meaning |- |} {| border="1" |+ I ! Acronym ! Meaning |- | IDCAF || interim Data Capture and Format |- | IPT || Integrated Product Team |- |} = J K L = {| border="1" |+ J ! Acronym ! Meaning |- |} {| border="1" |+ K ! Acronym ! Meaning |- |} {| border="1" |+ L ! Acronym ! Meaning |- | LCP || Left Circular Polarization |- | LNA || Low Noise Amplifier |- | LO || Local Oscillator |- |} = M N O = {| border="1" |+ M ! Acronym ! Meaning |- | MIB || Module Interface Board |- |} {| border="1" |+ N ! Acronym ! Meaning |- |} {| border="1" |+ O ! Acronym ! Meaning |- | OPT || Observation Preparation Tool |- | OTS || On-the sky |- | OMT || Ortho mode Transducer |- |} = P Q R = {| border="1" |+ P ! Acronym ! Meaning |- |} {| border="1" |+ Q ! Acronym ! Meaning |- |} {| border="1" |+ R ! Acronym ! Meaning |- | RFI || Radio Frequency Interference |- | RTCAT || Real-time Calibrator Analysis Tool |- |} = S T U = {| border="1" |+ S ! Acronym ! Meaning |- | SBC || Single Board Computer |- | SLC || Serial Line Controller |- |} {| border="1" |+ T ! Acronym ! Meaning |- |} {| border="1" |+ U ! Acronym ! Meaning |- |} = V W X = {| border="1" |+ V ! Acronym ! Meaning |- | VLA || Very Large Array |- | VLBA || Very Long Baseline Array |- | VME || Virtual Machine Environment |- | VPN || Virtual Private Network |- |} {| border="1" |+ W ! Acronym ! Meaning |- |} {| border="1" |+ X ! Acronym ! Meaning |- |} = Y Z = {| border="1" |+ Y ! Acronym ! Meaning |- |} {| border="1" |+ Z ! Acronym ! Meaning |- |} a860821286b1171499dc9966f5bb9d33bf07d8d0 133 131 2010-05-10T14:37:15Z Jmcmulli 2 wikitext text/x-wiki = A B C = {| border="1" |+ A ! Acronym ! Meaning |- | ACU || Antenna Control Unit |- | AOC || Array Operations Center (now DSOC) |- |} {| border="1" |+ B ! Acronym ! Meaning |- |} {| border="1" |+ C ! Acronym ! Meaning |- | CDL || Central Development Laboratory |- | CMP || Control and Monitor Processor |- |} = D E F = {| border="1" |+ D ! Acronym ! Meaning |- | DCAF || Data Capture and Format |- | DSOC || Domenici Science Operations Center |- | DTS || Digital Transmission System |- | FRM || Focus Rotation Mount |} {| border="1" |+ E ! Acronym ! Meaning |- | ECSO || EVLA Commissioning Staff Observing |} {| border="1" |+ F ! Acronym ! Meaning |- | FET || Field Effect Transistor |- |} = G H I = {| border="1" |+ G ! Acronym ! Meaning |- | GBT || Green Bank Telescope |- |} {| border="1" |+ H ! Acronym ! Meaning |- |} {| border="1" |+ I ! Acronym ! Meaning |- | IDCAF || interim Data Capture and Format |- | IPT || Integrated Product Team |- |} = J K L = {| border="1" |+ J ! Acronym ! Meaning |- |} {| border="1" |+ K ! Acronym ! Meaning |- |} {| border="1" |+ L ! Acronym ! Meaning |- | LCP || Left Circular Polarization |- | LNA || Low Noise Amplifier |- | LO || Local Oscillator |- |} = M N O = {| border="1" |+ M ! Acronym ! Meaning |- | MIB || Module Interface Board |- |} {| border="1" |+ N ! Acronym ! Meaning |- |} {| border="1" |+ O ! Acronym ! Meaning |- | OPT || Observation Preparation Tool |- | OSRO || Open Shared Risk Observing |- | OTS || On-the sky |- | OMT || Ortho mode Transducer |- |} = P Q R = {| border="1" |+ P ! Acronym ! Meaning |- |} {| border="1" |+ Q ! Acronym ! Meaning |- |} {| border="1" |+ R ! Acronym ! Meaning |- | RSRO || Resident Shared Risk Observing |- | RFI || Radio Frequency Interference |- | RTCAT || Real-time Calibrator Analysis Tool |- |} = S T U = {| border="1" |+ S ! Acronym ! Meaning |- | SBC || Single Board Computer |- | SLC || Serial Line Controller |- |} {| border="1" |+ T ! Acronym ! Meaning |- |} {| border="1" |+ U ! Acronym ! Meaning |- |} = V W X = {| border="1" |+ V ! Acronym ! Meaning |- | VLA || Very Large Array |- | VLBA || Very Long Baseline Array |- | VME || Virtual Machine Environment |- | VPN || Virtual Private Network |- |} {| border="1" |+ W ! Acronym ! Meaning |- |} {| border="1" |+ X ! Acronym ! Meaning |- |} = Y Z = {| border="1" |+ Y ! Acronym ! Meaning |- |} {| border="1" |+ Z ! Acronym ! Meaning |- |} 6bff16c39e865d02ea492aa2fa0345a159f9dd39 134 133 2010-05-10T20:34:26Z Jmcmulli 2 wikitext text/x-wiki = A B C = {| border="1" |+ A ! Acronym ! Meaning |- | ACU || Antenna Control Unit |- | AOC || Array Operations Center (now DSOC) |- |} {| border="1" |+ B ! Acronym ! Meaning |- |} {| border="1" |+ C ! Acronym ! Meaning |- | CDL || Central Development Laboratory |- | CMP || Control and Monitor Processor |- |} = D E F = {| border="1" |+ D ! Acronym ! Meaning |- | DCAF || Data Capture and Format |- | DSOC || Domenici Science Operations Center |- | DTS || Digital Transmission System |- | FPGA || Field Programmable Gate Array |- | FRM || Focus Rotation Mount |} {| border="1" |+ E ! Acronym ! Meaning |- | ECSO || EVLA Commissioning Staff Observing |} {| border="1" |+ F ! Acronym ! Meaning |- | FET || Field Effect Transistor |- |} = G H I = {| border="1" |+ G ! Acronym ! Meaning |- | GBT || Green Bank Telescope |- |} {| border="1" |+ H ! Acronym ! Meaning |- |} {| border="1" |+ I ! Acronym ! Meaning |- | IDCAF || interim Data Capture and Format |- | IPT || Integrated Product Team |- |} = J K L = {| border="1" |+ J ! Acronym ! Meaning |- |} {| border="1" |+ K ! Acronym ! Meaning |- |} {| border="1" |+ L ! Acronym ! Meaning |- | LCP || Left Circular Polarization |- | LNA || Low Noise Amplifier |- | LO || Local Oscillator |- |} = M N O = {| border="1" |+ M ! Acronym ! Meaning |- | MIB || Module Interface Board |- |} {| border="1" |+ N ! Acronym ! Meaning |- | NIC | Network Interface Card |- |} {| border="1" |+ O ! Acronym ! Meaning |- | OPT || Observation Preparation Tool |- | OSRO || Open Shared Risk Observing |- | OTS || On-the sky |- | OMT || Ortho mode Transducer |- |} = P Q R = {| border="1" |+ P ! Acronym ! Meaning |- |} {| border="1" |+ Q ! Acronym ! Meaning |- |} {| border="1" |+ R ! Acronym ! Meaning |- | RSRO || Resident Shared Risk Observing |- | RFI || Radio Frequency Interference |- | RTCAT || Real-time Calibrator Analysis Tool |- |} = S T U = {| border="1" |+ S ! Acronym ! Meaning |- | SBC || Single Board Computer |- | SLC || Serial Line Controller |- |} {| border="1" |+ T ! Acronym ! Meaning |- |} {| border="1" |+ U ! Acronym ! Meaning |- |} = V W X = {| border="1" |+ V ! Acronym ! Meaning |- | VLA || Very Large Array |- | VLBA || Very Long Baseline Array |- | VME || Virtual Machine Environment |- | VPN || Virtual Private Network |- |} {| border="1" |+ W ! Acronym ! Meaning |- | WIC || WIDAR Insanity Confirmation |- |} {| border="1" |+ X ! Acronym ! Meaning |- |} = Y Z = {| border="1" |+ Y ! Acronym ! Meaning |- |} {| border="1" |+ Z ! Acronym ! Meaning |- |} e28f5c6c8fe03af7f68004861709c849d19255e6 135 134 2010-05-10T20:35:11Z Jmcmulli 2 wikitext text/x-wiki = A B C = {| border="1" |+ A ! Acronym ! Meaning |- | ACU || Antenna Control Unit |- | AOC || Array Operations Center (now DSOC) |- |} {| border="1" |+ B ! Acronym ! Meaning |- |} {| border="1" |+ C ! Acronym ! Meaning |- | CDL || Central Development Laboratory |- | CMP || Control and Monitor Processor |- |} = D E F = {| border="1" |+ D ! Acronym ! Meaning |- | DCAF || Data Capture and Format |- | DSOC || Domenici Science Operations Center |- | DTS || Digital Transmission System |- | FPGA || Field Programmable Gate Array |- | FRM || Focus Rotation Mount |} {| border="1" |+ E ! Acronym ! Meaning |- | ECSO || EVLA Commissioning Staff Observing |} {| border="1" |+ F ! Acronym ! Meaning |- | FET || Field Effect Transistor |- |} = G H I = {| border="1" |+ G ! Acronym ! Meaning |- | GBT || Green Bank Telescope |- |} {| border="1" |+ H ! Acronym ! Meaning |- |} {| border="1" |+ I ! Acronym ! Meaning |- | IDCAF || interim Data Capture and Format |- | IPT || Integrated Product Team |- |} = J K L = {| border="1" |+ J ! Acronym ! Meaning |- |} {| border="1" |+ K ! Acronym ! Meaning |- |} {| border="1" |+ L ! Acronym ! Meaning |- | LCP || Left Circular Polarization |- | LNA || Low Noise Amplifier |- | LO || Local Oscillator |- |} = M N O = {| border="1" |+ M ! Acronym ! Meaning |- | MIB || Module Interface Board |- |} {| border="1" |+ N ! Acronym ! Meaning |- | NIC || Network Interface Card |- |} {| border="1" |+ O ! Acronym ! Meaning |- | OPT || Observation Preparation Tool |- | OSRO || Open Shared Risk Observing |- | OTS || On-the sky |- | OMT || Ortho mode Transducer |- |} = P Q R = {| border="1" |+ P ! Acronym ! Meaning |- |} {| border="1" |+ Q ! Acronym ! Meaning |- |} {| border="1" |+ R ! Acronym ! Meaning |- | RSRO || Resident Shared Risk Observing |- | RFI || Radio Frequency Interference |- | RTCAT || Real-time Calibrator Analysis Tool |- |} = S T U = {| border="1" |+ S ! Acronym ! Meaning |- | SBC || Single Board Computer |- | SLC || Serial Line Controller |- |} {| border="1" |+ T ! Acronym ! Meaning |- |} {| border="1" |+ U ! Acronym ! Meaning |- |} = V W X = {| border="1" |+ V ! Acronym ! Meaning |- | VLA || Very Large Array |- | VLBA || Very Long Baseline Array |- | VME || Virtual Machine Environment |- | VPN || Virtual Private Network |- |} {| border="1" |+ W ! Acronym ! Meaning |- | WIC || WIDAR Insanity Confirmation |- |} {| border="1" |+ X ! Acronym ! Meaning |- |} = Y Z = {| border="1" |+ Y ! Acronym ! Meaning |- |} {| border="1" |+ Z ! Acronym ! Meaning |- |} 5fa588b8c40f9ad4369f570a92ca63c102d83de5 142 135 2010-06-24T22:03:17Z Jmcmulli 2 wikitext text/x-wiki = A B C = {| border="1" |+ A ! Acronym ! Meaning |- | ACU || Antenna Control Unit |- | AOC || Array Operations Center (now DSOC) |- |} {| border="1" |+ B ! Acronym ! Meaning |- | BlB || Baseline Board |- |} {| border="1" |+ C ! Acronym ! Meaning |- | CBE || Correlator BackEnd |- | CDL || Central Development Laboratory |- | CM || Configuration Mapper |- | CMIB || Correlator Module Interface Board |- | CMP || Control and Monitor Processor |- | CPCC || Power Control Correlator Computer |- |} = D E F = {| border="1" |+ D ! Acronym ! Meaning |- | DCAF || Data Capture and Format |- | DDS || Direct Digital Synthesizer |- | DSOC || Domenici Science Operations Center |- | DTS || Digital Transmission System |- | FFT || Fast Fourier Transform |- | FPGA || Field Programmable Gate Array |- | FRM || Focus Rotation Mount |} {| border="1" |+ E ! Acronym ! Meaning |- | ECSO || EVLA Commissioning Staff Observing |} {| border="1" |+ F ! Acronym ! Meaning |- | FET || Field Effect Transistor |- |} = G H I = {| border="1" |+ G ! Acronym ! Meaning |- | GBT || Green Bank Telescope |- |} {| border="1" |+ H ! Acronym ! Meaning |- |} {| border="1" |+ I ! Acronym ! Meaning |- | IDCAF || interim Data Capture and Format |- | IPT || Integrated Product Team |- |} = J K L = {| border="1" |+ J ! Acronym ! Meaning |- |} {| border="1" |+ K ! Acronym ! Meaning |- |} {| border="1" |+ L ! Acronym ! Meaning |- | LCP || Left Circular Polarization |- | LNA || Low Noise Amplifier |- | LO || Local Oscillator |- |} = M N O = {| border="1" |+ M ! Acronym ! Meaning |- | MCAF || |- | MCCC || Master Correlator Control Computer |- | MIB || Module Interface Board |- |} {| border="1" |+ N ! Acronym ! Meaning |- | NIC || Network Interface Card |- |} {| border="1" |+ O ! Acronym ! Meaning |- | OPT || Observation Preparation Tool |- | OSRO || Open Shared Risk Observing |- | OTS || On-the sky |- | OMT || Ortho mode Transducer |- |} = P Q R = {| border="1" |+ P ! Acronym ! Meaning |- |} {| border="1" |+ Q ! Acronym ! Meaning |- |} {| border="1" |+ R ! Acronym ! Meaning |- | RSRO || Resident Shared Risk Observing |- | RFI || Radio Frequency Interference |- | RTCAT || Real-time Calibrator Analysis Tool |- |} = S T U = {| border="1" |+ S ! Acronym ! Meaning |- | SBC || Single Board Computer |- | SLC || Serial Line Controller |- | StB || Station Board |- |} {| border="1" |+ T ! Acronym ! Meaning |- | TelCal || Telescope Calibration |- |} {| border="1" |+ U ! Acronym ! Meaning |- |} = V W X = {| border="1" |+ V ! Acronym ! Meaning |- | VCI || Virtual Correlator Interface |- | VLA || Very Large Array |- | VLBA || Very Long Baseline Array |- | VME || Virtual Machine Environment |- | VPN || Virtual Private Network |- |} {| border="1" |+ W ! Acronym ! Meaning |- | WIC || WIDAR Insanity Confirmation |- |} {| border="1" |+ X ! Acronym ! Meaning |- | XML || eXtensible Markup Language |- |} = Y Z = {| border="1" |+ Y ! Acronym ! Meaning |- | YIG || Yttrium Iron Garnet |- |} {| border="1" |+ Z ! Acronym ! Meaning |- |} 699c9fac126e94012c271cb495229595d9405b2e 146 142 2010-06-28T22:24:45Z Jmcmulli 2 wikitext text/x-wiki = A B C = {| border="1" |+ A ! Acronym ! Meaning |- | ACU || Antenna Control Unit |- | ALC || Automatic Level Control |- | AOC || Array Operations Center (now DSOC) |- |} {| border="1" |+ B ! Acronym ! Meaning |- | BlB || Baseline Board |- |} {| border="1" |+ C ! Acronym ! Meaning |- | CBE || Correlator BackEnd |- | CDL || Central Development Laboratory |- | CM || Configuration Mapper |- | CMIB || Correlator Module Interface Board |- | CMP || Control and Monitor Processor |- | CPCC || Power Control Correlator Computer |- |} = D E F = {| border="1" |+ D ! Acronym ! Meaning |- | DCAF || Data Capture and Format |- | DDS || Direct Digital Synthesizer |- | DSOC || Domenici Science Operations Center |- | DTS || Digital Transmission System |- | FFT || Fast Fourier Transform |- | FPGA || Field Programmable Gate Array |- | FRM || Focus Rotation Mount |} {| border="1" |+ E ! Acronym ! Meaning |- | ECSO || EVLA Commissioning Staff Observing |} {| border="1" |+ F ! Acronym ! Meaning |- | FET || Field Effect Transistor |- |} = G H I = {| border="1" |+ G ! Acronym ! Meaning |- | GBT || Green Bank Telescope |- |} {| border="1" |+ H ! Acronym ! Meaning |- |} {| border="1" |+ I ! Acronym ! Meaning |- | IDCAF || interim Data Capture and Format |- | IPT || Integrated Product Team |- |} = J K L = {| border="1" |+ J ! Acronym ! Meaning |- |} {| border="1" |+ K ! Acronym ! Meaning |- |} {| border="1" |+ L ! Acronym ! Meaning |- | LCP || Left Circular Polarization |- | LNA || Low Noise Amplifier |- | LO || Local Oscillator |- |} = M N O = {| border="1" |+ M ! Acronym ! Meaning |- | MCAF || |- | MCCC || Master Correlator Control Computer |- | MIB || Module Interface Board |- |} {| border="1" |+ N ! Acronym ! Meaning |- | NIC || Network Interface Card |- |} {| border="1" |+ O ! Acronym ! Meaning |- | OPT || Observation Preparation Tool |- | OSRO || Open Shared Risk Observing |- | OTS || On-the sky |- | OMT || Ortho mode Transducer |- |} = P Q R = {| border="1" |+ P ! Acronym ! Meaning |- |} {| border="1" |+ Q ! Acronym ! Meaning |- |} {| border="1" |+ R ! Acronym ! Meaning |- | RSRO || Resident Shared Risk Observing |- | RFI || Radio Frequency Interference |- | RTCAT || Real-time Calibrator Analysis Tool |- |} = S T U = {| border="1" |+ S ! Acronym ! Meaning |- | SBC || Single Board Computer |- | SLC || Serial Line Controller |- | StB || Station Board |- |} {| border="1" |+ T ! Acronym ! Meaning |- | TelCal || Telescope Calibration |- |} {| border="1" |+ U ! Acronym ! Meaning |- |} = V W X = {| border="1" |+ V ! Acronym ! Meaning |- | VCI || Virtual Correlator Interface |- | VLA || Very Large Array |- | VLBA || Very Long Baseline Array |- | VME || Virtual Machine Environment |- | VPN || Virtual Private Network |- |} {| border="1" |+ W ! Acronym ! Meaning |- | WIC || WIDAR Insanity Confirmation |- |} {| border="1" |+ X ! Acronym ! Meaning |- | XML || eXtensible Markup Language |- |} = Y Z = {| border="1" |+ Y ! Acronym ! Meaning |- | YIG || Yttrium Iron Garnet |- |} {| border="1" |+ Z ! Acronym ! Meaning |- |} 1f6d2e591fe5afa4bd9c73a8a7494d38b367fee9 0 1 132 48 2010-05-07T20:12:57Z Jmcmulli 2 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">Featured Article</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{Featured Article}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">News</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | * 12-Apr-2010: First 27-antenna correlation {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:left;" | <div style="font-size:120%">Configuration & Proposal Dates</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| |- | style="padding-left:6px; padding-top:6px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| ! Trimester ! Observing Period ! Configuration ! Proposal Deadline |- | 2010a | 2010 Mar 01 - 2010 Jul 12 | align="center" | D | 2009 Oct 1 |- | 2010a | 2010 Jul 16 - 2010 Aug 02 | align="center" | DnC | 2009 Oct 1 |- | 2010b | 2010 Aug 13 - 2010 Nov 01 | align="center" | C | 2010 Feb 1 |- | 2010b | 2010 Nov 05 - 2010 Nov 22 | align="center" | CnB | 2010 Feb 1 |- | 2010c | 2010 Dec 03 - 2011 Feb 22 | align="center" | B | 2010 Jun 1 |- | 2010c | 2011 Feb 25 - 2011 Mar 14 | align="center" | BnA | 2010 Jun 1 |- | 2011a | 2011 Mar 25 - 2011 Jun 27 | align='center' | A | 2010 Oct 1 |} <!-- TEXT 2 (END) --> |} |} 03484a0c3e74d631bfbd5d9f6b368ebced1942b7 143 132 2010-06-28T00:31:50Z Jmcmulli 2 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">Featured Article</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{Featured Article}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">News</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | * 02-Mar-2010: First science observations with WIDAR * 12-Apr-2010: First 27-antenna correlation {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:left;" | <div style="font-size:120%">Configuration & Proposal Dates</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| |- | style="padding-left:6px; padding-top:6px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| ! Trimester ! Observing Period ! Configuration ! Proposal Deadline |- | 2010a | 2010 Mar 01 - 2010 Jul 12 | align="center" | D | 2009 Oct 1 |- | 2010a | 2010 Jul 16 - 2010 Aug 02 | align="center" | DnC | 2009 Oct 1 |- | 2010b | 2010 Aug 13 - 2010 Nov 01 | align="center" | C | 2010 Feb 1 |- | 2010b | 2010 Nov 05 - 2010 Nov 22 | align="center" | CnB | 2010 Feb 1 |- | 2010c | 2010 Dec 03 - 2011 Feb 22 | align="center" | B | 2010 Jun 1 |- | 2010c | 2011 Feb 25 - 2011 Mar 14 | align="center" | BnA | 2010 Jun 1 |- | 2011a | 2011 Mar 25 - 2011 Jun 27 | align='center' | A | 2010 Oct 1 |} <!-- TEXT 2 (END) --> |} |} 79740968916b27e6b8cf6618e803e16ebc7086b0 145 143 2010-06-28T14:07:08Z Jmcmulli 2 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">Featured Article</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{Featured Article}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">News</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | * 02-Mar-2010: First science observations with WIDAR * 12-Apr-2010: First 27-antenna correlation * 25-Jun-2010: First fringes with 3-bit samplers (3 antennas) {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:left;" | <div style="font-size:120%">Configuration & Proposal Dates</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| |- | style="padding-left:6px; padding-top:6px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| ! Trimester ! Observing Period ! Configuration ! Proposal Deadline |- | 2010a | 2010 Mar 01 - 2010 Jul 12 | align="center" | D | 2009 Oct 1 |- | 2010a | 2010 Jul 16 - 2010 Aug 02 | align="center" | DnC | 2009 Oct 1 |- | 2010b | 2010 Aug 13 - 2010 Nov 01 | align="center" | C | 2010 Feb 1 |- | 2010b | 2010 Nov 05 - 2010 Nov 22 | align="center" | CnB | 2010 Feb 1 |- | 2010c | 2010 Dec 03 - 2011 Feb 22 | align="center" | B | 2010 Jun 1 |- | 2010c | 2011 Feb 25 - 2011 Mar 14 | align="center" | BnA | 2010 Jun 1 |- | 2011a | 2011 Mar 25 - 2011 Jun 27 | align='center' | A | 2010 Oct 1 |} <!-- TEXT 2 (END) --> |} |} 9956cf77738dc045911f492544e3e7f78ab69712 6 17 138 2010-05-12T07:08:21Z Cchandle 6 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 0 18 149 2010-07-19T20:54:39Z Jmcmulli 2 Created page with '_calcodes_ are generated based on observing intents established during the SB creation. These intents are then passed along to the SDM data set. _Note: These will eventually be r…' wikitext text/x-wiki _calcodes_ are generated based on observing intents established during the SB creation. These intents are then passed along to the SDM data set. _Note: These will eventually be replaced by scan intents. These are also summarized currently in AIPS by typing:_ <pre> >help calcode </pre> Giving: <pre> Help on CALCODE in AIPS version 31DEC10 CALCODE Type: Adverb (String*4) Use: This string is used to specify a desired calibration source code. ' ' => any calibrator code selected '* ' => any non blank code (calibrators only) '-CAL' => blank codes only (no calibrators) Null value: none The VLA has attached meaning to letters as follows: None = target source, or no potential calcode specified Old style VLA (indicators of positional accuracy): A = positional accuracy <0.002 arcseconds B = positional accuracy 0.002 - 0.01 arcseconds C = positional accuracy 0.01 - 0.15 arcseconds T = positional accuracy >0.15 arcseconds New style EVLA (indicators of scan intents): D = calibrator useful to determine Complex Gains E = calibrator useful to determine the absolute Flux Density Scale; typically only used for 3C48, 3C147, 3C286 (and 3C138, 3C295) F = calibrator useful to determine the spectral Bandpass response G = calibrator useful to determine Polarization Angle H = calibrator useful to determine the instrumental Polarization Leakage I = calibrator for Complex Gain and Bandpass J = calibrator for Complex Gain and Polarization Leakage K = calibrator for Flux Density Scale and Bandpass L = calibrator for Flux Density Scale and Polarization Angle M = calibrator for Bandpass and Polarization Angle N = calibrator for Flux Density Scale, Bandpass and Polarization Angle P = calibrator for Pointing observations Y = recognised as calibrator source in the VLA catalog by the pipeline; only appears when no other calcodes are present and NOT an indication that this calibrator is useful at this frequency or in this array Z = any other non-Target combination of intents Occasional VLBA: V = recognised as VLBI source in the VLBI "sched" catalog Code Gain Flux BP Pang P% D X E X F X G X H X I X X AIPS 1: ** press RETURN for more, enter Q or next line to quit print ** # autarch AIPS (31DEC10) 2492 06-MAY-2010 14:50:41 Page 2 Help on CALCODE in AIPS version 31DEC10 J X X K X X L X X M X X N X X X The VLBA has attached different meanings on occasion as well to some of the EVLA codes. b35f40005e639be611c86a67799bfc755c3af43b 150 149 2010-07-19T20:55:14Z Jmcmulli 2 wikitext text/x-wiki _calcodes_ are generated based on observing intents established during the SB creation. These intents are then passed along to the SDM data set. _Note: These will eventually be replaced by scan intents. These are also summarized currently in AIPS by typing:_ <pre> >help calcode </pre> Giving: <pre> Help on CALCODE in AIPS version 31DEC10 CALCODE Type: Adverb (String*4) Use: This string is used to specify a desired calibration source code. ' ' => any calibrator code selected '* ' => any non blank code (calibrators only) '-CAL' => blank codes only (no calibrators) Null value: none The VLA has attached meaning to letters as follows: None = target source, or no potential calcode specified Old style VLA (indicators of positional accuracy): A = positional accuracy <0.002 arcseconds B = positional accuracy 0.002 - 0.01 arcseconds C = positional accuracy 0.01 - 0.15 arcseconds T = positional accuracy >0.15 arcseconds New style EVLA (indicators of scan intents): D = calibrator useful to determine Complex Gains E = calibrator useful to determine the absolute Flux Density Scale; typically only used for 3C48, 3C147, 3C286 (and 3C138, 3C295) F = calibrator useful to determine the spectral Bandpass response G = calibrator useful to determine Polarization Angle H = calibrator useful to determine the instrumental Polarization Leakage I = calibrator for Complex Gain and Bandpass J = calibrator for Complex Gain and Polarization Leakage K = calibrator for Flux Density Scale and Bandpass L = calibrator for Flux Density Scale and Polarization Angle M = calibrator for Bandpass and Polarization Angle N = calibrator for Flux Density Scale, Bandpass and Polarization Angle P = calibrator for Pointing observations Y = recognised as calibrator source in the VLA catalog by the pipeline; only appears when no other calcodes are present and NOT an indication that this calibrator is useful at this frequency or in this array Z = any other non-Target combination of intents Occasional VLBA: V = recognised as VLBI source in the VLBI "sched" catalog Code Gain Flux BP Pang P% D X E X F X G X H X I X X J X X K X X L X X M X X N X X X The VLBA has attached different meanings on occasion as well to some of the EVLA codes. b22fbc7fb902563c1eedec40fa5f295c24988b99 0 1 151 145 2010-08-06T13:59:04Z Jmcmulli 2 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">Featured Article</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{Featured Article}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">News</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | * 02-Mar-2010: First science observations with WIDAR * 12-Apr-2010: First 27-antenna correlation * 25-Jun-2010: First fringes with 3-bit samplers (3 antennas) {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:left;" | <div style="font-size:120%">Configuration & Proposal Dates</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| |- | style="padding-left:6px; padding-top:6px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| ! Trimester ! Observing Period ! Configuration ! Proposal Deadline |- | 2010 Mar 01 - 2010 Sep 13 | align="center" | D | 2009 Oct 1 |- | 2010 Sep 17 - 2010 Oct 04 | align="center" | DnC | 2009 Oct 1 |- | 2010 Oct 15 - 2011 Jan 03 | align="center" | C | 2010 Feb 1 |- | 2011 Jan 07 - 2011 Jan 24 | align="center" | CnB | 2010 Feb 1 |- | 2011 Feb 04 - 2011 Apr 25 | align="center" | B | 2010 Jun 1 |- | 2011 Apr 29 - 2011 May 16 | align="center" | BnA | 2010 Jun 1 |- | 2011 May 27 - 2011 Aug 29 | align='center' | A | 2010 Oct 1 |} <!-- TEXT 2 (END) --> |} |} bfa2e105d703e51ddcc637db481cb59652f20529 0 16 152 146 2010-08-16T19:32:08Z Jmcmulli 2 /* V W X */ wikitext text/x-wiki = A B C = {| border="1" |+ A ! Acronym ! Meaning |- | ACU || Antenna Control Unit |- | ALC || Automatic Level Control |- | AOC || Array Operations Center (now DSOC) |- |} {| border="1" |+ B ! Acronym ! Meaning |- | BlB || Baseline Board |- |} {| border="1" |+ C ! Acronym ! Meaning |- | CBE || Correlator BackEnd |- | CDL || Central Development Laboratory |- | CM || Configuration Mapper |- | CMIB || Correlator Module Interface Board |- | CMP || Control and Monitor Processor |- | CPCC || Power Control Correlator Computer |- |} = D E F = {| border="1" |+ D ! Acronym ! Meaning |- | DCAF || Data Capture and Format |- | DDS || Direct Digital Synthesizer |- | DSOC || Domenici Science Operations Center |- | DTS || Digital Transmission System |- | FFT || Fast Fourier Transform |- | FPGA || Field Programmable Gate Array |- | FRM || Focus Rotation Mount |} {| border="1" |+ E ! Acronym ! Meaning |- | ECSO || EVLA Commissioning Staff Observing |} {| border="1" |+ F ! Acronym ! Meaning |- | FET || Field Effect Transistor |- |} = G H I = {| border="1" |+ G ! Acronym ! Meaning |- | GBT || Green Bank Telescope |- |} {| border="1" |+ H ! Acronym ! Meaning |- |} {| border="1" |+ I ! Acronym ! Meaning |- | IDCAF || interim Data Capture and Format |- | IPT || Integrated Product Team |- |} = J K L = {| border="1" |+ J ! Acronym ! Meaning |- |} {| border="1" |+ K ! Acronym ! Meaning |- |} {| border="1" |+ L ! Acronym ! Meaning |- | LCP || Left Circular Polarization |- | LNA || Low Noise Amplifier |- | LO || Local Oscillator |- |} = M N O = {| border="1" |+ M ! Acronym ! Meaning |- | MCAF || |- | MCCC || Master Correlator Control Computer |- | MIB || Module Interface Board |- |} {| border="1" |+ N ! Acronym ! Meaning |- | NIC || Network Interface Card |- |} {| border="1" |+ O ! Acronym ! Meaning |- | OPT || Observation Preparation Tool |- | OSRO || Open Shared Risk Observing |- | OTS || On-the sky |- | OMT || Ortho mode Transducer |- |} = P Q R = {| border="1" |+ P ! Acronym ! Meaning |- |} {| border="1" |+ Q ! Acronym ! Meaning |- |} {| border="1" |+ R ! Acronym ! Meaning |- | RSRO || Resident Shared Risk Observing |- | RFI || Radio Frequency Interference |- | RTCAT || Real-time Calibrator Analysis Tool |- |} = S T U = {| border="1" |+ S ! Acronym ! Meaning |- | SBC || Single Board Computer |- | SLC || Serial Line Controller |- | StB || Station Board |- |} {| border="1" |+ T ! Acronym ! Meaning |- | TelCal || Telescope Calibration |- |} {| border="1" |+ U ! Acronym ! Meaning |- |} = V W X = {| border="1" |+ V ! Acronym ! Meaning |- | VCI || Virtual Correlator Interface |- | VLA || Very Large Array |- | VLBA || Very Long Baseline Array |- | VME || Virtual Machine Environment |- | VPN || Virtual Private Network |- |} {| border="1" |+ W ! Acronym ! Meaning |- | WIC || WIDAR Insanity Confirmation |- | WIDAR || Wideband Interferometric Digital ARchitecture |- |} {| border="1" |+ X ! Acronym ! Meaning |- | XML || eXtensible Markup Language |- |} = Y Z = {| border="1" |+ Y ! Acronym ! Meaning |- | YIG || Yttrium Iron Garnet |- |} {| border="1" |+ Z ! Acronym ! Meaning |- |} 8d2706e5b90ad9b028b7e28c16373e5a4e00436f 153 152 2010-08-16T19:33:04Z Jmcmulli 2 /* M N O */ wikitext text/x-wiki = A B C = {| border="1" |+ A ! Acronym ! Meaning |- | ACU || Antenna Control Unit |- | ALC || Automatic Level Control |- | AOC || Array Operations Center (now DSOC) |- |} {| border="1" |+ B ! Acronym ! Meaning |- | BlB || Baseline Board |- |} {| border="1" |+ C ! Acronym ! Meaning |- | CBE || Correlator BackEnd |- | CDL || Central Development Laboratory |- | CM || Configuration Mapper |- | CMIB || Correlator Module Interface Board |- | CMP || Control and Monitor Processor |- | CPCC || Power Control Correlator Computer |- |} = D E F = {| border="1" |+ D ! Acronym ! Meaning |- | DCAF || Data Capture and Format |- | DDS || Direct Digital Synthesizer |- | DSOC || Domenici Science Operations Center |- | DTS || Digital Transmission System |- | FFT || Fast Fourier Transform |- | FPGA || Field Programmable Gate Array |- | FRM || Focus Rotation Mount |} {| border="1" |+ E ! Acronym ! Meaning |- | ECSO || EVLA Commissioning Staff Observing |} {| border="1" |+ F ! Acronym ! Meaning |- | FET || Field Effect Transistor |- |} = G H I = {| border="1" |+ G ! Acronym ! Meaning |- | GBT || Green Bank Telescope |- |} {| border="1" |+ H ! Acronym ! Meaning |- |} {| border="1" |+ I ! Acronym ! Meaning |- | IDCAF || interim Data Capture and Format |- | IPT || Integrated Product Team |- |} = J K L = {| border="1" |+ J ! Acronym ! Meaning |- |} {| border="1" |+ K ! Acronym ! Meaning |- |} {| border="1" |+ L ! Acronym ! Meaning |- | LCP || Left Circular Polarization |- | LNA || Low Noise Amplifier |- | LO || Local Oscillator |- |} = M N O = {| border="1" |+ M ! Acronym ! Meaning |- | MCAF || Metadata Capture and Formatting |- | MCCC || Master Correlator Control Computer |- | MIB || Module Interface Board |- |} {| border="1" |+ N ! Acronym ! Meaning |- | NIC || Network Interface Card |- |} {| border="1" |+ O ! Acronym ! Meaning |- | OPT || Observation Preparation Tool |- | OSRO || Open Shared Risk Observing |- | OTS || On-the sky |- | OMT || Ortho mode Transducer |- |} = P Q R = {| border="1" |+ P ! Acronym ! Meaning |- |} {| border="1" |+ Q ! Acronym ! Meaning |- |} {| border="1" |+ R ! Acronym ! Meaning |- | RSRO || Resident Shared Risk Observing |- | RFI || Radio Frequency Interference |- | RTCAT || Real-time Calibrator Analysis Tool |- |} = S T U = {| border="1" |+ S ! Acronym ! Meaning |- | SBC || Single Board Computer |- | SLC || Serial Line Controller |- | StB || Station Board |- |} {| border="1" |+ T ! Acronym ! Meaning |- | TelCal || Telescope Calibration |- |} {| border="1" |+ U ! Acronym ! Meaning |- |} = V W X = {| border="1" |+ V ! Acronym ! Meaning |- | VCI || Virtual Correlator Interface |- | VLA || Very Large Array |- | VLBA || Very Long Baseline Array |- | VME || Virtual Machine Environment |- | VPN || Virtual Private Network |- |} {| border="1" |+ W ! Acronym ! Meaning |- | WIC || WIDAR Insanity Confirmation |- | WIDAR || Wideband Interferometric Digital ARchitecture |- |} {| border="1" |+ X ! Acronym ! Meaning |- | XML || eXtensible Markup Language |- |} = Y Z = {| border="1" |+ Y ! Acronym ! Meaning |- | YIG || Yttrium Iron Garnet |- |} {| border="1" |+ Z ! Acronym ! Meaning |- |} 7ffaf0b940288e818ebecb5836b4926bc2f65797 154 153 2010-08-16T19:34:24Z Jmcmulli 2 /* A B C */ wikitext text/x-wiki = A B C = {| border="1" |+ A ! Acronym ! Meaning |- | ACU || Antenna Control Unit |- | ALC || Automatic Level Control |- | AOC || Array Operations Center (now DSOC) |- |} {| border="1" |+ B ! Acronym ! Meaning |- | BlB || Baseline Board |- |} {| border="1" |+ C ! Acronym ! Meaning |- | CBE || Correlator BackEnd |- | CDL || Central Development Laboratory |- | CM || Configuration Mapper |- | CMIB || Correlator Module Interface Board |- | CMP || Control and Monitor Processor |- | CPCC || Power Control Correlator Computer |- | CRC || Cyclic Redundancy Check |- |} = D E F = {| border="1" |+ D ! Acronym ! Meaning |- | DCAF || Data Capture and Format |- | DDS || Direct Digital Synthesizer |- | DSOC || Domenici Science Operations Center |- | DTS || Digital Transmission System |- | FFT || Fast Fourier Transform |- | FPGA || Field Programmable Gate Array |- | FRM || Focus Rotation Mount |} {| border="1" |+ E ! Acronym ! Meaning |- | ECSO || EVLA Commissioning Staff Observing |} {| border="1" |+ F ! Acronym ! Meaning |- | FET || Field Effect Transistor |- |} = G H I = {| border="1" |+ G ! Acronym ! Meaning |- | GBT || Green Bank Telescope |- |} {| border="1" |+ H ! Acronym ! Meaning |- |} {| border="1" |+ I ! Acronym ! Meaning |- | IDCAF || interim Data Capture and Format |- | IPT || Integrated Product Team |- |} = J K L = {| border="1" |+ J ! Acronym ! Meaning |- |} {| border="1" |+ K ! Acronym ! Meaning |- |} {| border="1" |+ L ! Acronym ! Meaning |- | LCP || Left Circular Polarization |- | LNA || Low Noise Amplifier |- | LO || Local Oscillator |- |} = M N O = {| border="1" |+ M ! Acronym ! Meaning |- | MCAF || Metadata Capture and Formatting |- | MCCC || Master Correlator Control Computer |- | MIB || Module Interface Board |- |} {| border="1" |+ N ! Acronym ! Meaning |- | NIC || Network Interface Card |- |} {| border="1" |+ O ! Acronym ! Meaning |- | OPT || Observation Preparation Tool |- | OSRO || Open Shared Risk Observing |- | OTS || On-the sky |- | OMT || Ortho mode Transducer |- |} = P Q R = {| border="1" |+ P ! Acronym ! Meaning |- |} {| border="1" |+ Q ! Acronym ! Meaning |- |} {| border="1" |+ R ! Acronym ! Meaning |- | RSRO || Resident Shared Risk Observing |- | RFI || Radio Frequency Interference |- | RTCAT || Real-time Calibrator Analysis Tool |- |} = S T U = {| border="1" |+ S ! Acronym ! Meaning |- | SBC || Single Board Computer |- | SLC || Serial Line Controller |- | StB || Station Board |- |} {| border="1" |+ T ! Acronym ! Meaning |- | TelCal || Telescope Calibration |- |} {| border="1" |+ U ! Acronym ! Meaning |- |} = V W X = {| border="1" |+ V ! Acronym ! Meaning |- | VCI || Virtual Correlator Interface |- | VLA || Very Large Array |- | VLBA || Very Long Baseline Array |- | VME || Virtual Machine Environment |- | VPN || Virtual Private Network |- |} {| border="1" |+ W ! Acronym ! Meaning |- | WIC || WIDAR Insanity Confirmation |- | WIDAR || Wideband Interferometric Digital ARchitecture |- |} {| border="1" |+ X ! Acronym ! Meaning |- | XML || eXtensible Markup Language |- |} = Y Z = {| border="1" |+ Y ! Acronym ! Meaning |- | YIG || Yttrium Iron Garnet |- |} {| border="1" |+ Z ! Acronym ! Meaning |- |} 52e2397a29eca0833cad6681fb107c9fb075f1fc 156 154 2010-08-31T21:16:35Z Jmcmulli 2 wikitext text/x-wiki = A B C = {| border="1" |+ A ! Acronym ! Meaning |- | ACU || Antenna Control Unit |- | ALC || Automatic Level Control |- | AOC || Array Operations Center (now DSOC) |- |} {| border="1" |+ B ! Acronym ! Meaning |- | BlB || Baseline Board |- |} {| border="1" |+ C ! Acronym ! Meaning |- | CASA || Common Astronomy Software Applications |- | CBE || Correlator BackEnd |- | CDL || Central Development Laboratory |- | CM || Configuration Mapper |- | CMIB || Correlator Module Interface Board |- | CMP || Control and Monitor Processor |- | CPCC || Power Control Correlator Computer |- | CRC || Cyclic Redundancy Check |- |} = D E F = {| border="1" |+ D ! Acronym ! Meaning |- | DCAF || Data Capture and Format |- | DDS || Direct Digital Synthesizer |- | DiFX || Distributed FX software correlator used for VLBA |- | DSOC || Domenici Science Operations Center |- | DTS || Digital Transmission System |- | FFT || Fast Fourier Transform |- | FPGA || Field Programmable Gate Array |- | FRM || Focus Rotation Mount |} {| border="1" |+ E ! Acronym ! Meaning |- | ECSO || EVLA Commissioning Staff Observing |} {| border="1" |+ F ! Acronym ! Meaning |- | FET || Field Effect Transistor |- |} = G H I = {| border="1" |+ G ! Acronym ! Meaning |- | GBT || Green Bank Telescope |- |} {| border="1" |+ H ! Acronym ! Meaning |- |} {| border="1" |+ I ! Acronym ! Meaning |- | IDCAF || interim Data Capture and Format |- | IPT || Integrated Product Team |- |} = J K L = {| border="1" |+ J ! Acronym ! Meaning |- |} {| border="1" |+ K ! Acronym ! Meaning |- |} {| border="1" |+ L ! Acronym ! Meaning |- | LCP || Left Circular Polarization |- | LNA || Low Noise Amplifier |- | LO || Local Oscillator |- |} = M N O = {| border="1" |+ M ! Acronym ! Meaning |- | MCAF || Metadata Capture and Formatting |- | MCCC || Master Correlator Control Computer |- | MIB || Module Interface Board |- |} {| border="1" |+ N ! Acronym ! Meaning |- | NGAS || Next Generation Archive Server |- | NIC || Network Interface Card |- |} {| border="1" |+ O ! Acronym ! Meaning |- | OPT || Observation Preparation Tool |- | OSRO || Open Shared Risk Observing |- | OTS || On-the sky |- | OMT || Ortho mode Transducer |- |} = P Q R = {| border="1" |+ P ! Acronym ! Meaning |- | PST || Proposal Submission Tool |- |} {| border="1" |+ Q ! Acronym ! Meaning |- |} {| border="1" |+ R ! Acronym ! Meaning |- | RSRO || Resident Shared Risk Observing |- | RFI || Radio Frequency Interference |- | RTCAT || Real-time Calibrator Analysis Tool |- | RXP || Retiming, Crossbar & Phasing |- |} = S T U = {| border="1" |+ S ! Acronym ! Meaning |- | SBC || Single Board Computer |- | SDM || Science Data Model |- | SEFD || System Equivalent Flux Density |- | SLC || Serial Line Controller |- | StB || Station Board |- |} {| border="1" |+ T ! Acronym ! Meaning |- | TelCal || Telescope Calibration |- |} {| border="1" |+ U ! Acronym ! Meaning |- |} = V W X = {| border="1" |+ V ! Acronym ! Meaning |- | VCI || Virtual Correlator Interface |- | VEX || VLBI EXchange (observation description file format) |- | VLA || Very Large Array |- | VLBA || Very Long Baseline Array |- | VLBI || Very Long Baseline Interferometry |- | VME || Virtual Machine Environment |- | VPN || Virtual Private Network |- |} {| border="1" |+ W ! Acronym ! Meaning |- | WIC || WIDAR Insanity Confirmation |- | WIDAR || Wideband Interferometric Digital ARchitecture |- |} {| border="1" |+ X ! Acronym ! Meaning |- | XML || eXtensible Markup Language |- |} = Y Z = {| border="1" |+ Y ! Acronym ! Meaning |- | YIG || Yttrium Iron Garnet |- |} {| border="1" |+ Z ! Acronym ! Meaning |- |} a30ed3ef32dd55698995d176c483146c6471de28 174 156 2010-10-27T21:02:09Z Jmcmulli 2 wikitext text/x-wiki = A B C = {| border="1" |+ A ! Acronym ! Meaning |- | ACU || Antenna Control Unit |- | ALC || Automatic Level Control |- | AOC || Array Operations Center (now DSOC) |- |} {| border="1" |+ B ! Acronym ! Meaning |- | BlB || Baseline Board |- |} {| border="1" |+ C ! Acronym ! Meaning |- | CASA || Common Astronomy Software Applications |- | CBE || Correlator BackEnd |- | CDL || Central Development Laboratory |- | CM || Configuration Mapper |- | CMIB || Correlator Module Interface Board |- | CMP || Control and Monitor Processor |- | CPCC || Power Control Correlator Computer |- | CRC || Cyclic Redundancy Check |- |} = D E F = {| border="1" |+ D ! Acronym ! Meaning |- | DCAF || Data Capture and Format |- | DDS || Direct Digital Synthesizer |- | DiFX || Distributed FX software correlator used for VLBA |- | DMTV || Delay Module Test Vector |- | DSOC || Domenici Science Operations Center |- | DTS || Digital Transmission System |- | FFT || Fast Fourier Transform |- | FPGA || Field Programmable Gate Array |- | FRM || Focus Rotation Mount |} {| border="1" |+ E ! Acronym ! Meaning |- | ECSO || EVLA Commissioning Staff Observing |} {| border="1" |+ F ! Acronym ! Meaning |- | FET || Field Effect Transistor |- |} = G H I = {| border="1" |+ G ! Acronym ! Meaning |- | GBT || Green Bank Telescope |- |} {| border="1" |+ H ! Acronym ! Meaning |- |} {| border="1" |+ I ! Acronym ! Meaning |- | IDCAF || interim Data Capture and Format |- | IPT || Integrated Product Team |- |} = J K L = {| border="1" |+ J ! Acronym ! Meaning |- |} {| border="1" |+ K ! Acronym ! Meaning |- |} {| border="1" |+ L ! Acronym ! Meaning |- | LCP || Left Circular Polarization |- | LNA || Low Noise Amplifier |- | LO || Local Oscillator |- |} = M N O = {| border="1" |+ M ! Acronym ! Meaning |- | MCAF || Metadata Capture and Formatting |- | MCCC || Master Correlator Control Computer |- | MIB || Module Interface Board |- |} {| border="1" |+ N ! Acronym ! Meaning |- | NGAS || Next Generation Archive Server |- | NIC || Network Interface Card |- |} {| border="1" |+ O ! Acronym ! Meaning |- | OPT || Observation Preparation Tool |- | OSRO || Open Shared Risk Observing |- | OTS || On-the sky |- | OMT || Ortho mode Transducer |- |} = P Q R = {| border="1" |+ P ! Acronym ! Meaning |- | PST || Proposal Submission Tool |- |} {| border="1" |+ Q ! Acronym ! Meaning |- |} {| border="1" |+ R ! Acronym ! Meaning |- | RSRO || Resident Shared Risk Observing |- | RFI || Radio Frequency Interference |- | RTCAT || Real-time Calibrator Analysis Tool |- | RXP || Retiming, Crossbar & Phasing |- |} = S T U = {| border="1" |+ S ! Acronym ! Meaning |- | SBC || Single Board Computer |- | SDM || Science Data Model |- | SEFD || System Equivalent Flux Density |- | SLC || Serial Line Controller |- | StB || Station Board |- |} {| border="1" |+ T ! Acronym ! Meaning |- | TelCal || Telescope Calibration |- |} {| border="1" |+ U ! Acronym ! Meaning |- |} = V W X = {| border="1" |+ V ! Acronym ! Meaning |- | VCI || Virtual Correlator Interface |- | VEX || VLBI EXchange (observation description file format) |- | VLA || Very Large Array |- | VLBA || Very Long Baseline Array |- | VLBI || Very Long Baseline Interferometry |- | VME || Virtual Machine Environment |- | VPN || Virtual Private Network |- |} {| border="1" |+ W ! Acronym ! Meaning |- | WIC || WIDAR Insanity Confirmation |- | WIDAR || Wideband Interferometric Digital ARchitecture |- |} {| border="1" |+ X ! Acronym ! Meaning |- | XML || eXtensible Markup Language |- |} = Y Z = {| border="1" |+ Y ! Acronym ! Meaning |- | YIG || Yttrium Iron Garnet |- |} {| border="1" |+ Z ! Acronym ! Meaning |- |} 8c1c56399cbb7163a01a08a8527a9ee41dba8106 10 2 155 148 2010-08-26T17:43:30Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [http://mctest.evla.nrao.edu/cgi-bin/evla/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [http://mctest.evla.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [http://mctest.evla.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] · [[Key to Calcodes]] |} beb35c9a2199949e63ff16b613a7cd15ad7450b4 173 155 2010-09-28T14:25:48Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] · [[Key to Calcodes]] |} 68a136232b89bae5bd5f998aba1ee9ff9bab774a 175 173 2010-10-28T21:26:10Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:Observing Preparation| Spectral Line Observations]] · [[:Category:Observing Preparation| Polarimetry Observations]] · [[:Category:Observing Preparation| Planetary Observations]] · [[:Category:Observing Preparation| High Frequency Observing (K, Ka, Q)]] · [[:Category:Observing Preparation| Low Frequency Observing (L, S, C)]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] · [[Key to Calcodes]] |} 4101b720d946598861930a6c938aca8ecff2c1cf 194 175 2010-12-02T17:07:18Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:SpectraLine| Spectral Line Observations]] · [[:Category:Polarimetry| Polarimetry Observations]] · [[:Category:Planetary| Planetary Observations]] · [[:Category:HighFrequency| High Frequency Observing (K, Ka, Q)]] · [[:Category:LowFrequency| Low Frequency Observing (L, S, C)]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] · [[Key to Calcodes]] |} d1ca9a7e749d6b2fca2bc67c01d1372fc0920eff 200 194 2010-12-10T15:20:00Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] · [[Observational Status Summary - Projected]] [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:SpectraLine| Spectral Line Observations]] · [[:Category:Polarimetry| Polarimetry Observations]] · [[:Category:Planetary| Planetary Observations]] · [[:Category:HighFrequency| High Frequency Observing (K, Ka, Q)]] · [[:Category:LowFrequency| Low Frequency Observing (L, S, C)]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] · [[Key to Calcodes]] |} 134ef45eed292dd9d9ce2343737b18da5705486f 14 8 157 140 2010-09-01T21:28:38Z Jmcmulli 2 /* An Overview of the EVLA */ wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 12, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q3 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Sep ! Oct-Jan ! Feb-Apr |- | 2010 || '''D''' || '''C''' || '''B''' |- ! Year ! May-Aug ! Sep-Dec ! Jan-Apr | 2011 || '''A''' || '''D''' || '''C''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim we plan to mount and test the compatibility of the existing 74-MHz dipoles with the wideband EVLA electronics in the upcoming C-configuration, with the goal of providing a low frequency observing capability in the B/BnA/A configurations. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, July 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 8 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 23 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 1 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 4 || align='center'| 4 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 24 || align='center'| 24 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|none|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 60ddcaffe36ed7868006479d657828843f09d003 158 157 2010-09-01T21:33:21Z Jmcmulli 2 /* An Overview of the EVLA */ wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 12, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q3 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' || - |- ! Year ! Feb-Apr ! May-Aug ! Sep-Dec |- | 2011 || '''B''' || '''A''' || '''D''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim we plan to mount and test the compatibility of the existing 74-MHz dipoles with the wideband EVLA electronics in the upcoming C-configuration, with the goal of providing a low frequency observing capability in the B/BnA/A configurations. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, July 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 8 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 23 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 1 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 4 || align='center'| 4 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 24 || align='center'| 24 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|none|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. b8822e49df9574bcd0b4894933438f15382829df 159 158 2010-09-01T21:39:57Z Jmcmulli 2 /* Resolution */ wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 12, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q3 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' || - |- ! Year ! Feb-Apr ! May-Aug ! Sep-Dec |- | 2011 || '''B''' || '''A''' || '''D''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim we plan to mount and test the compatibility of the existing 74-MHz dipoles with the wideband EVLA electronics in the upcoming C-configuration, with the goal of providing a low frequency observing capability in the B/BnA/A configurations. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, July 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 8 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 23 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 1 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 4 || align='center'| 4 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 24 || align='center'| 24 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24.0 || 80.0 || 260.0 || 850.0 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800.0 || 2200.0 || 20000.0 || 20000.0 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|none|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. be940685adfba42ec6bdb0a5371b7e7e59c6e79a 160 159 2010-09-01T21:40:52Z Jmcmulli 2 /* Resolution */ wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 12, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q3 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' || - |- ! Year ! Feb-Apr ! May-Aug ! Sep-Dec |- | 2011 || '''B''' || '''A''' || '''D''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim we plan to mount and test the compatibility of the existing 74-MHz dipoles with the wideband EVLA electronics in the upcoming C-configuration, with the goal of providing a low frequency observing capability in the B/BnA/A configurations. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, July 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 8 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 23 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 1 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 4 || align='center'| 4 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 24 || align='center'| 24 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24 || 80 || 260 || 850 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800 || 2200 || 20000 || 20000 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|none|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 15610f1445d7bd67b7f42ce64bd536830d2ce565 161 160 2010-09-01T21:42:47Z Jmcmulli 2 /* Sensitivity */ wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 12, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q3 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' || - |- ! Year ! Feb-Apr ! May-Aug ! Sep-Dec |- | 2011 || '''B''' || '''A''' || '''D''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim we plan to mount and test the compatibility of the existing 74-MHz dipoles with the wideband EVLA electronics in the upcoming C-configuration, with the goal of providing a low frequency observing capability in the B/BnA/A configurations. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, July 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 8 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 23 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 1 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 4 || align='center'| 4 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 24 || align='center'| 24 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24 || 80 || 260 || 850 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800 || 2200 || 20000 || 20000 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|none|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. No current estimates are available for 4-band (74 MHz) observations but will be established during October 2010 DnC configuration testing. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 65364a966c852d2504d7c0f893f8701010071f04 162 161 2010-09-01T22:09:51Z Jmcmulli 2 /* An Overview of the EVLA */ wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 12, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q3 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' |- ! Year ! Feb-Apr ! May-Aug ! Sep-Dec |- | 2011 || '''B''' || '''A''' || '''D''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim we plan to mount and test the compatibility of the existing 74-MHz dipoles with the wideband EVLA electronics in the upcoming C-configuration, with the goal of providing a low frequency observing capability in the B/BnA/A configurations. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, July 2010 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 8 || align='center'| 26 |- | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10 |- | 6 cm (C) || 4.0-8.0 || align='center'| 23 || align='center'| 26 |- | 3 cm (X) || 8.0-12.0 || align='center'| 1 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 4 || align='center'| 4 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 24 || align='center'| 24 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24 || 80 || 260 || 850 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800 || 2200 || 20000 || 20000 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|none|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. No current estimates are available for 4-band (74 MHz) observations but will be established during October 2010 DnC configuration testing. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 4e4be462519a28fd47e73bff138d8142fc96a7fc 164 162 2010-09-01T22:42:14Z Jmcmulli 2 /* Expected Capabilities */ wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 12, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q3 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' |- ! Year ! Feb-Apr ! May-Aug ! Sep-Dec |- | 2011 || '''B''' || '''A''' || '''D''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim we plan to mount and test the compatibility of the existing 74-MHz dipoles with the wideband EVLA electronics in the upcoming C-configuration, with the goal of providing a low frequency observing capability in the B/BnA/A configurations. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in May 2011, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, May 2011 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 16 || align='center'| 27 |- | 13 cm (S) || 2.0-4.0 || align='center'| 16 || align='center'| 16 |- | 6 cm (C) || 4.0-8.0 || align='center'| 27 || align='center'| 27 |- | 3 cm (X) || 8.0-12.0 || align='center'| 7 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 11 || align='center'| 11 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 27 || align='center'| 27 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24 || 80 || 260 || 850 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800 || 2200 || 20000 || 20000 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|none|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. No current estimates are available for 4-band (74 MHz) observations but will be established during October 2010 DnC configuration testing. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. d075997e3094464ef324f7e3ecfbf137288162ee 165 164 2010-09-01T22:42:50Z Jmcmulli 2 /* What is the Expanded Very Large Array? */ wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 12, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q4 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' |- ! Year ! Feb-Apr ! May-Aug ! Sep-Dec |- | 2011 || '''B''' || '''A''' || '''D''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim we plan to mount and test the compatibility of the existing 74-MHz dipoles with the wideband EVLA electronics in the upcoming C-configuration, with the goal of providing a low frequency observing capability in the B/BnA/A configurations. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in May 2011, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, May 2011 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 16 || align='center'| 27 |- | 13 cm (S) || 2.0-4.0 || align='center'| 16 || align='center'| 16 |- | 6 cm (C) || 4.0-8.0 || align='center'| 27 || align='center'| 27 |- | 3 cm (X) || 8.0-12.0 || align='center'| 7 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 11 || align='center'| 11 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 27 || align='center'| 27 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24 || 80 || 260 || 850 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800 || 2200 || 20000 || 20000 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|none|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. No current estimates are available for 4-band (74 MHz) observations but will be established during October 2010 DnC configuration testing. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 403140678862403da81e3a8519e1d308bdae26c6 166 165 2010-09-01T22:49:42Z Jmcmulli 2 /* An Overview of the EVLA */ wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 12, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q4 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' |- ! Year ! Feb-Apr ! May-Aug ! Sep-Dec |- | 2011 || '''B''' || '''A''' || '''D''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim we plan to mount and test the compatibility of a set of six 74-MHz dipoles with the wideband EVLA electronics in the upcoming C-configuration, with the goal of providing a low frequency observing capability (approximately 4.5 MHz BW in the 68-86 MHz range) in the B/BnA/A configurations; the sensitivity should be assumed to be that of the old VLA system (RMS (10 min)=160 mJy), though the improved correlator, bandwidth and interference environment may significantly improve this value. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in May 2011, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, May 2011 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 16 || align='center'| 27 |- | 13 cm (S) || 2.0-4.0 || align='center'| 16 || align='center'| 16 |- | 6 cm (C) || 4.0-8.0 || align='center'| 27 || align='center'| 27 |- | 3 cm (X) || 8.0-12.0 || align='center'| 7 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 11 || align='center'| 11 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 27 || align='center'| 27 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24 || 80 || 260 || 850 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800 || 2200 || 20000 || 20000 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|none|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. No current estimates are available for 4-band (74 MHz) observations but will be established during October 2010 DnC configuration testing. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 6ed918e3e66ba307a4b98932ad20d4e53e2ab2bc 167 166 2010-09-22T17:53:53Z Jmcmulli 2 /* An Overview of the EVLA */ wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 12, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q4 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' |- ! Year ! Feb-Apr ! May-Aug ! Sep-Dec |- | 2011 || '''B''' || '''A''' || '''D''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim, a set of six 74-MHz dipoles were mounted with the wideband EVLA electronics in the DnC-configuration, providing information on the low frequency observing capability (response over 16 MHz BW in the 62-78 MHz range, with a significant roll-off in the low frequency end due to the dipole response); the October 1 2010 call for proposals will accept observations in this band for the B/BnA/A configurations; the sensitivity should be assumed to be that of the old VLA system (RMS (10 min)=160 mJy), though the improved correlator, bandwidth and interference environment may significantly improve this value. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in May 2011, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, May 2011 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 16 || align='center'| 27 |- | 13 cm (S) || 2.0-4.0 || align='center'| 16 || align='center'| 16 |- | 6 cm (C) || 4.0-8.0 || align='center'| 27 || align='center'| 27 |- | 3 cm (X) || 8.0-12.0 || align='center'| 7 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 11 || align='center'| 11 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 27 || align='center'| 27 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24 || 80 || 260 || 850 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800 || 2200 || 20000 || 20000 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|none|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. No current estimates are available for 4-band (74 MHz) observations but will be established during October 2010 DnC configuration testing. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. b78269b50e38f589ddf4fb020d936a50529dce57 169 167 2010-09-22T18:02:37Z Jmcmulli 2 /* An Overview of the EVLA */ wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 12, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q4 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' |- ! Year ! Feb-Apr ! May-Aug ! Sep-Dec |- | 2011 || '''B''' || '''A''' || '''D''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim, a set of six 74-MHz dipoles were mounted with the wideband EVLA electronics in the DnC-configuration, providing information on the low frequency observing capability (response over 16 MHz BW in the 62-78 MHz range, with a significant roll-off in the low frequency end due to the dipole response); the October 1 2010 call for proposals will accept observations in this band for the B/BnA/A configurations; the sensitivity should be assumed to be that of the old VLA system (RMS (10 min)=160 mJy), though the improved correlator, bandwidth and interference environment may significantly improve this value (See: http://evlaguides.nrao.edu/index.php?title=File:4-bandResponse.png). The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in May 2011, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, May 2011 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 20 cm (L) || 1.0-2.0 || align='center'| 16 || align='center'| 27 |- | 13 cm (S) || 2.0-4.0 || align='center'| 16 || align='center'| 16 |- | 6 cm (C) || 4.0-8.0 || align='center'| 27 || align='center'| 27 |- | 3 cm (X) || 8.0-12.0 || align='center'| 7 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 11 || align='center'| 11 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 27 || align='center'| 27 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24 || 80 || 260 || 850 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800 || 2200 || 20000 || 20000 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|none|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. No current estimates are available for 4-band (74 MHz) observations but will be established during October 2010 DnC configuration testing. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 55798b6e85c860b919cd6e8b61d398e74421b3a3 170 169 2010-09-23T17:37:07Z Jmcmulli 2 /* Expected Capabilities */ wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 12, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q4 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' |- ! Year ! Feb-Apr ! May-Aug ! Sep-Dec |- | 2011 || '''B''' || '''A''' || '''D''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim, a set of six 74-MHz dipoles were mounted with the wideband EVLA electronics in the DnC-configuration, providing information on the low frequency observing capability (response over 16 MHz BW in the 62-78 MHz range, with a significant roll-off in the low frequency end due to the dipole response); the October 1 2010 call for proposals will accept observations in this band for the B/BnA/A configurations; the sensitivity should be assumed to be that of the old VLA system (RMS (10 min)=160 mJy), though the improved correlator, bandwidth and interference environment may significantly improve this value (See: http://evlaguides.nrao.edu/index.php?title=File:4-bandResponse.png). The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in May 2011, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, May 2011 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 400 cm (4-band) || .062 - 0.078 || || align='center'| 27 |- | 20 cm (L) || 1.0-2.0 || align='center'| 16 || align='center'| 27 |- | 13 cm (S) || 2.0-4.0 || align='center'| 16 || align='center'| 16 |- | 6 cm (C) || 4.0-8.0 || align='center'| 27 || align='center'| 27 |- | 3 cm (X) || 8.0-12.0 || align='center'| 7 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 11 || align='center'| 11 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 27 || align='center'| 27 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24 || 80 || 260 || 850 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800 || 2200 || 20000 || 20000 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|none|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. No current estimates are available for 4-band (74 MHz) observations but will be established during October 2010 DnC configuration testing. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 63f36a007fa4cc6a5c61d7ec321b0d30a2b57f88 171 170 2010-09-23T17:37:44Z Jmcmulli 2 /* Expected Capabilities */ wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 12, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q4 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' |- ! Year ! Feb-Apr ! May-Aug ! Sep-Dec |- | 2011 || '''B''' || '''A''' || '''D''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim, a set of six 74-MHz dipoles were mounted with the wideband EVLA electronics in the DnC-configuration, providing information on the low frequency observing capability (response over 16 MHz BW in the 62-78 MHz range, with a significant roll-off in the low frequency end due to the dipole response); the October 1 2010 call for proposals will accept observations in this band for the B/BnA/A configurations; the sensitivity should be assumed to be that of the old VLA system (RMS (10 min)=160 mJy), though the improved correlator, bandwidth and interference environment may significantly improve this value (See: http://evlaguides.nrao.edu/index.php?title=File:4-bandResponse.png). The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in May 2011, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, May 2011 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 400 cm (4-band) || .062 - .078 || || align='center'| 27 |- | 20 cm (L) || 1.0-2.0 || align='center'| 16 || align='center'| 27 |- | 13 cm (S) || 2.0-4.0 || align='center'| 16 || align='center'| 16 |- | 6 cm (C) || 4.0-8.0 || align='center'| 27 || align='center'| 27 |- | 3 cm (X) || 8.0-12.0 || align='center'| 7 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 11 || align='center'| 11 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 27 || align='center'| 27 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24 || 80 || 260 || 850 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800 || 2200 || 20000 || 20000 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|none|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. No current estimates are available for 4-band (74 MHz) observations but will be established during October 2010 DnC configuration testing. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 479013c14da2745156998ba485617ddffadb1c30 172 171 2010-09-23T17:38:05Z Jmcmulli 2 /* Expected Capabilities */ wikitext text/x-wiki '''The EVLA Observational Status Summary''' ''Version date: May 12, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q4 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' |- ! Year ! Feb-Apr ! May-Aug ! Sep-Dec |- | 2011 || '''B''' || '''A''' || '''D''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim, a set of six 74-MHz dipoles were mounted with the wideband EVLA electronics in the DnC-configuration, providing information on the low frequency observing capability (response over 16 MHz BW in the 62-78 MHz range, with a significant roll-off in the low frequency end due to the dipole response); the October 1 2010 call for proposals will accept observations in this band for the B/BnA/A configurations; the sensitivity should be assumed to be that of the old VLA system (RMS (10 min)=160 mJy), though the improved correlator, bandwidth and interference environment may significantly improve this value (See: http://evlaguides.nrao.edu/index.php?title=File:4-bandResponse.png). The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in May 2011, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, May 2011 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 400 cm (4-band) || 0.062-0.078 || || align='center'| 27 |- | 20 cm (L) || 1.0-2.0 || align='center'| 16 || align='center'| 27 |- | 13 cm (S) || 2.0-4.0 || align='center'| 16 || align='center'| 16 |- | 6 cm (C) || 4.0-8.0 || align='center'| 27 || align='center'| 27 |- | 3 cm (X) || 8.0-12.0 || align='center'| 7 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 11 || align='center'| 11 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 27 || align='center'| 27 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24 || 80 || 260 || 850 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800 || 2200 || 20000 || 20000 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|none|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. No current estimates are available for 4-band (74 MHz) observations but will be established during October 2010 DnC configuration testing. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 76e476a5c2e33462a89906b24a41826b2051e1e0 6 17 163 138 2010-09-01T22:35:23Z Jmcmulli 2 uploaded a new version of &quot;[[File:WideBandRcvrFrcstMay10.png]]&quot;: Updated to August 2010. wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 19 168 2010-09-22T17:58:39Z Jmcmulli 2 Plot of the frequency response of Cygnus A observed with WIDAR over a 16 MHz bandpass on one representative baseline. The response peaks at 74 MHz but has signal over 16 MHz; at the upper under the response is limited by a ~14 MHz filter centered at 70 MH wikitext text/x-wiki Plot of the frequency response of Cygnus A observed with WIDAR over a 16 MHz bandpass on one representative baseline. The response peaks at 74 MHz but has signal over 16 MHz; at the upper under the response is limited by a ~14 MHz filter centered at 70 MHz, while the lower end rolls off due to the dipole response. 4d010f1a4b96781dd33ba25c1ba2ad9861602d92 14 20 176 2010-11-09T15:27:06Z Jmcmulli 2 Created page with "= Polarization Calibration = == Current OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observati..." wikitext text/x-wiki = Polarization Calibration = == Current OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Revised Guidelines == === Observing Guidelines === Three strategies for deriving the instrumental leakage terms: * Single scan observation of a zero polarization source * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source * Two scans of a source of known polarization ==== Low Frequency Considerations ==== * TEC, Ionosphere Solar activity Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== ==== Frequency stability ==== ==== Polarization Calibrator Catalog ==== Zero pol High pol ==== Monitoring Observations ==== === Post-processing Guidelines === 509896d396fe01c2fa10e7072b73c0e3329b929e 177 176 2010-11-09T16:00:46Z Jmcmulli 2 wikitext text/x-wiki = Polarization Calibration = == Current OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Revised Guidelines == === Observing Recommendations === There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (see catalog below) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source * Two scans of a source of known polarization (see catalog below) ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog ==== CSO (Compact Symmetric Objects) are characteristically unpolarized and can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78); these make up the bulk of the zero/unpolarized sources. {| border="1" align="center" |+ '''Table 1: Unpolarized sources''' !Source !RA (1950) !DEC (1950) |- | 0026+346 || 00 26 34.8386 || 34 39 57.586 |- | 0108+388 || 01 08 47.2595 || 38 50 32.691 |- | 0134+329 (3C48) || 01 34 49.8264 || 32 54 20.259 |- | 0316+413 (3C84) || 03 16 29.5673 || 41 19 51.916 |- | 0404+768 || 04 03 58.60 || 76 48 54.0 |- | 0710+439 || 07 10 03.3460 || 43 54 26.216 |- | 1031+567 || 10 31 55.9562 || 56 44 18.284 |- | 1146+596 || 11 46 10.4160 || 59 41 36.834 |- | 1358+624 || 13 58 58.310 || 62 25 08.40 |- | 1404+286 (OQ208) || 14 04 45.6151 || 28 41 29.235 |- | 1815+614 || 18 15 05.4851 || 61 26 04.496 |- | 1826+796 || 18 26 43.2676 || 79 36 59.943 |- | 1943+546 || 19 43 22.6729 || 54 40 47.955 |- | 1946+708 || 19 46 12.0492 || 70 48 21.397 |- | 2021+614 || 20 21 13.3005 || 61 27 18.157 |- | 2352+495 || 23 52 37.7919 || 49 33 26.701 |- |} High pol 3C148 3C286 ==== Monitoring Observations ==== Description and scope. List of observation dates and results. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. 3334a2274e5ae859867a542763b28010943c55fe 178 177 2010-11-16T20:00:35Z Smyers 4 /* Current OSS Guidelines */ wikitext text/x-wiki = Polarization Calibration = == Current OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Revised Guidelines == === Observing Recommendations === There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (see catalog below) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source * Two scans of a source of known polarization (see catalog below) ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog ==== CSO (Compact Symmetric Objects) are characteristically unpolarized and can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78); these make up the bulk of the zero/unpolarized sources. {| border="1" align="center" |+ '''Table 1: Unpolarized sources''' !Source !RA (1950) !DEC (1950) |- | 0026+346 || 00 26 34.8386 || 34 39 57.586 |- | 0108+388 || 01 08 47.2595 || 38 50 32.691 |- | 0134+329 (3C48) || 01 34 49.8264 || 32 54 20.259 |- | 0316+413 (3C84) || 03 16 29.5673 || 41 19 51.916 |- | 0404+768 || 04 03 58.60 || 76 48 54.0 |- | 0710+439 || 07 10 03.3460 || 43 54 26.216 |- | 1031+567 || 10 31 55.9562 || 56 44 18.284 |- | 1146+596 || 11 46 10.4160 || 59 41 36.834 |- | 1358+624 || 13 58 58.310 || 62 25 08.40 |- | 1404+286 (OQ208) || 14 04 45.6151 || 28 41 29.235 |- | 1815+614 || 18 15 05.4851 || 61 26 04.496 |- | 1826+796 || 18 26 43.2676 || 79 36 59.943 |- | 1943+546 || 19 43 22.6729 || 54 40 47.955 |- | 1946+708 || 19 46 12.0492 || 70 48 21.397 |- | 2021+614 || 20 21 13.3005 || 61 27 18.157 |- | 2352+495 || 23 52 37.7919 || 49 33 26.701 |- |} High pol 3C148 3C286 ==== Monitoring Observations ==== Description and scope. List of observation dates and results. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. b7b18aafbf9b3089e261838df39f03ca45005d91 179 178 2010-11-16T20:03:39Z Smyers 4 /* Polarization Calibrator Catalog */ wikitext text/x-wiki = Polarization Calibration = == Current OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Revised Guidelines == === Observing Recommendations === There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (see catalog below) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source * Two scans of a source of known polarization (see catalog below) ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog ==== CSO (Compact Symmetric Objects) are characteristically unpolarized and can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78); these make up the bulk of the zero/unpolarized sources. WARNING: the source names given below are the B1950 names, we are in the process of updating this list to J2000 and the names used in the EVLA OPT. {| border="1" align="center" |+ '''Table 1: Unpolarized sources''' !Source !RA (1950) !DEC (1950) |- | 0026+346 || 00 26 34.8386 || 34 39 57.586 |- | 0108+388 || 01 08 47.2595 || 38 50 32.691 |- | 0134+329 (3C48) || 01 34 49.8264 || 32 54 20.259 |- | 0316+413 (3C84) || 03 16 29.5673 || 41 19 51.916 |- | 0404+768 || 04 03 58.60 || 76 48 54.0 |- | 0710+439 || 07 10 03.3460 || 43 54 26.216 |- | 1031+567 || 10 31 55.9562 || 56 44 18.284 |- | 1146+596 || 11 46 10.4160 || 59 41 36.834 |- | 1358+624 || 13 58 58.310 || 62 25 08.40 |- | 1404+286 (OQ208) || 14 04 45.6151 || 28 41 29.235 |- | 1815+614 || 18 15 05.4851 || 61 26 04.496 |- | 1826+796 || 18 26 43.2676 || 79 36 59.943 |- | 1943+546 || 19 43 22.6729 || 54 40 47.955 |- | 1946+708 || 19 46 12.0492 || 70 48 21.397 |- | 2021+614 || 20 21 13.3005 || 61 27 18.157 |- | 2352+495 || 23 52 37.7919 || 49 33 26.701 |- |} High pol 3C148 3C286 ==== Monitoring Observations ==== Description and scope. List of observation dates and results. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. 210ed4fb481b4e4e962c888a7157596d43ed125c 180 179 2010-11-16T23:15:05Z Smyers 4 /* Polarization Calibrator Catalog */ wikitext text/x-wiki = Polarization Calibration = == Current OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Revised Guidelines == === Observing Recommendations === There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (see catalog below) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source * Two scans of a source of known polarization (see catalog below) ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog ==== The following sources are known to be CSO (Compact Symmetric Objects) are characteristically unpolarized and can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators noted below. CSOs tend to be on the weak side and should be used with care at higher frequencies. There are also a few "bright, low pol" sources available. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 1: Verified unpolarized and polarized calibrator sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0137+3309 || 01 34 49.8264 || 32 54 20.259 || 0134+329 (3C48) || pol standard (<6cm) |- | J0319+4130 || 03 16 29.5673 || 41 19 51.916 || 0316+413 (3C84) || bright, low pol, flat spectrum |- | J0521+1638 || 05 18 16.5141 || 16 35 26.834 || 0518+165 (3C138) || pol standard |- | J0713+4349 || 07 10 03.3460 || 43 54 26.216 || 0710+439 || CSO, monitored |- | J1331+3030 || 13 28 49.6577 || 30 45 58.640 || 1328+307 (3C286) || pol standard |- | J1407+2827 || 14 04 45.6151 || 28 41 29.235 || 1404+286 (OQ208) || low pol, steep spectrum |- | J2355+4950 || 23 52 37.7919 || 49 33 26.701 || 2352+495 || CSO, monitored |- |} Comments: at least one "pol standard" should be included for angle calibration, "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ {| border="1" align="center" |+ '''Table 2: Unverified low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: northern CSO sources from Gugliucci et al. study (see above). ==== Monitoring Observations ==== Description and scope. List of observation dates and results. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. 8ae5752e2102ed64c5be634fa7d28b8a10186fd1 181 180 2010-11-16T23:24:56Z Smyers 4 /* Polarization Calibrator Catalog */ wikitext text/x-wiki = Polarization Calibration = == Current OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Revised Guidelines == === Observing Recommendations === There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (see catalog below) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source * Two scans of a source of known polarization (see catalog below) ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of bright sources with variable flux density and polarization (monitored by EVLA). There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). WARNING: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. {| border="1" align="center" |+ '''Table 1: Verified unpolarized and polarized calibrator sources''' !Source !Other name !Comments |- | J0137+3309 || B0134+329 (3C48) || pol standard (<6cm) |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum |- | J0521+1638 || B0518+165 (3C138) || pol standard |- | J0713+4349 || B0710+439 || CSO, monitored |- | J1331+3030 || B1328+307 (3C286) || pol standard |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum |- | J2355+4950 || B2352+495 || CSO, monitored |- |} Comments: * at least one "pol standard" should be included for angle calibration * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies The following sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 2: Unverified low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * northern CSO sources from Gugliucci et al. study (see above). ==== Monitoring Observations ==== Description and scope. List of observation dates and results. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. d0dc1435f56d2627871496f16b744f0806b6c3f9 182 181 2010-11-16T23:55:50Z Smyers 4 /* Polarization Calibrator Catalog */ wikitext text/x-wiki = Polarization Calibration = == Current OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Revised Guidelines == === Observing Recommendations === There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (see catalog below) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source * Two scans of a source of known polarization (see catalog below) ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. {| border="1" align="center" |+ '''Table 1: Verified unpolarized and polarized calibrator sources''' !Source !Other name !Comments |- | J0137+3309 || B0134+329 (3C48) || pol standard (<6cm) |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability |- | J0521+1638 || B0518+165 (3C138) || pol standard |- | J0555+3948 || B0552+398 (3C138) || bright, flat spectrum, monitored, moderate variability |- | J0713+4349 || B0710+439 || low pol, CSO, monitored |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability |- | J1331+3030 || B1328+307 (3C286) || pol standard |- | J1310+3220 || B1308+326 || monitored |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored |- | J2355+4950 || B2352+495 || low pol, CSO, monitored |- |} Comments: * at least one "pol standard" should be included for angle calibration * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies The following sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 2: Unverified low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * northern CSO sources from Gugliucci et al. study (see above). ==== Monitoring Observations ==== Description and scope. List of observation dates and results. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. 04971b8d7b26b59cf3c3eac9e1a8e87ff37b8183 183 182 2010-11-17T00:00:43Z Smyers 4 /* Monitoring Observations */ wikitext text/x-wiki = Polarization Calibration = == Current OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Revised Guidelines == === Observing Recommendations === There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (see catalog below) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source * Two scans of a source of known polarization (see catalog below) ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. {| border="1" align="center" |+ '''Table 1: Verified unpolarized and polarized calibrator sources''' !Source !Other name !Comments |- | J0137+3309 || B0134+329 (3C48) || pol standard (<6cm) |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability |- | J0521+1638 || B0518+165 (3C138) || pol standard |- | J0555+3948 || B0552+398 (3C138) || bright, flat spectrum, monitored, moderate variability |- | J0713+4349 || B0710+439 || low pol, CSO, monitored |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability |- | J1331+3030 || B1328+307 (3C286) || pol standard |- | J1310+3220 || B1308+326 || monitored |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored |- | J2355+4950 || B2352+495 || low pol, CSO, monitored |- |} Comments: * at least one "pol standard" should be included for angle calibration * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies The following sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 2: Unverified low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * northern CSO sources from Gugliucci et al. study (see above). ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. 2074cb673321c740e94ab1052851ac03463dcfdf 184 183 2010-11-17T00:03:44Z Smyers 4 /* Polarization Calibrator Catalog */ wikitext text/x-wiki = Polarization Calibration = == Current OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Revised Guidelines == === Observing Recommendations === There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (see catalog below) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source * Two scans of a source of known polarization (see catalog below) ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. {| border="1" align="center" |+ '''Table 1: Verified unpolarized and polarized calibrator sources''' !Source !Other name !Comments |- | J0137+3309 || B0134+329 (3C48) || pol standard (<6cm) |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability |- | J0521+1638 || B0518+165 (3C138) || pol standard |- | J0555+3948 || B0552+398 (3C138) || bright, flat spectrum, monitored, moderate variability |- | J0713+4349 || B0710+439 || low pol, CSO, monitored |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability |- | J1331+3030 || B1328+307 (3C286) || pol standard |- | J1310+3220 || B1308+326 || monitored |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored |- | J2355+4950 || B2352+495 || low pol, CSO, monitored |- |} Comments: * at least one "pol standard" should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) The following sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 2: Unverified low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * northern CSO sources from Gugliucci et al. study (see above). ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. fb71a28eb072f04ce85dbd604a957e8e2cb29956 185 184 2010-11-17T16:51:11Z Smyers 4 /* Polarization Calibrator Catalog */ wikitext text/x-wiki = Polarization Calibration = == Current OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Revised Guidelines == === Observing Recommendations === There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (see catalog below) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source * Two scans of a source of known polarization (see catalog below) ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standard sources''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (<6cm) || 1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || 2 |- |} Notes: *1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0555+3948 || B0552+398 (3C138) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || 3 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || 3 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || 3 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || 4 |- |} Notes: *3. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *4. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || 5 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || 6 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || 7 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || 6 |- |} Notes: *5. Very bright and low polarization above 4GHz (but variable flux density). *6. Weak at high frequency, but stable flux and very low polarization. *7. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. 78505649370edfa3565bd0a0c6f50c629aca4e12 186 185 2010-11-17T16:51:30Z Smyers 4 /* Polarization Calibrator Catalog and Selection */ wikitext text/x-wiki = Polarization Calibration = == Current OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Revised Guidelines == === Observing Recommendations === There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (see catalog below) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source * Two scans of a source of known polarization (see catalog below) ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (<6cm) || 1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || 2 |- |} Notes: *1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0555+3948 || B0552+398 (3C138) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || 3 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || 3 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || 3 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || 4 |- |} Notes: *3. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *4. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || 5 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || 6 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || 7 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || 6 |- |} Notes: *5. Very bright and low polarization above 4GHz (but variable flux density). *6. Weak at high frequency, but stable flux and very low polarization. *7. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. 457d9fc4ff27f4c7089123a22396aa77ba3879a4 187 186 2010-11-17T17:05:19Z Smyers 4 /* Current OSS Guidelines */ wikitext text/x-wiki = Polarization Calibration = == Current OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Revised Guidelines == === Observing Recommendations === There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (see catalog below) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source * Two scans of a source of known polarization (see catalog below) ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (<6cm) || 1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || 2 |- |} Notes: *1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0555+3948 || B0552+398 (3C138) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || 3 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || 3 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || 3 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || 4 |- |} Notes: *3. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *4. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || 5 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || 6 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || 7 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || 6 |- |} Notes: *5. Very bright and low polarization above 4GHz (but variable flux density). *6. Weak at high frequency, but stable flux and very low polarization. *7. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. 2773252e5af86b60b3d6fb62b3e4171c26ee72ca 188 187 2010-11-17T17:12:13Z Smyers 4 /* Observing Recommendations */ wikitext text/x-wiki = Polarization Calibration = == Current OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Revised Guidelines == === Observing Recommendations === There are several strategies for deriving the Q/U angle calibration: * Observation of a primary polarization standard (Category A) * Observation of a secondary polarization calibrator (Category B with Note 3) with auxilary monitoring observations to transfer from primary. This calibration is needed to set the polarization vector angle 0.5*arctan(U/Q) and should be done in all cases. There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (Category C) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source (Category B) * Two scans of a source of known polarization (Category A or B with transfer) See Tables 1-4 below for Category A-D source catalogs. ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (<6cm) || 1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || 2 |- |} Notes: *1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0555+3948 || B0552+398 (3C138) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || 3 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || 3 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || 3 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || 4 |- |} Notes: *3. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *4. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || 5 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || 6 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || 7 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || 6 |- |} Notes: *5. Very bright and low polarization above 4GHz (but variable flux density). *6. Weak at high frequency, but stable flux and very low polarization. *7. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. b93aec645dae6494ffde774198a2200714603d09 189 188 2010-11-17T17:27:35Z Jmcmulli 2 /* Current OSS Guidelines */ wikitext text/x-wiki = Polarization Calibration = == Revised OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Revised Guidelines == === Observing Recommendations === There are several strategies for deriving the Q/U angle calibration: * Observation of a primary polarization standard (Category A) * Observation of a secondary polarization calibrator (Category B with Note 3) with auxilary monitoring observations to transfer from primary. This calibration is needed to set the polarization vector angle 0.5*arctan(U/Q) and should be done in all cases. There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (Category C) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source (Category B) * Two scans of a source of known polarization (Category A or B with transfer) See Tables 1-4 below for Category A-D source catalogs. ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (<6cm) || 1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || 2 |- |} Notes: *1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0555+3948 || B0552+398 (3C138) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || 3 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || 3 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || 3 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || 4 |- |} Notes: *3. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *4. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || 5 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || 6 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || 7 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || 6 |- |} Notes: *5. Very bright and low polarization above 4GHz (but variable flux density). *6. Weak at high frequency, but stable flux and very low polarization. *7. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. cb5a18a98e6bd5a3443413f14335cb3960aa811d 190 189 2010-11-17T17:28:06Z Jmcmulli 2 /* Revised Guidelines */ wikitext text/x-wiki = Polarization Calibration = == Revised OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Detailed Guidelines == === Observing Recommendations === There are several strategies for deriving the Q/U angle calibration: * Observation of a primary polarization standard (Category A) * Observation of a secondary polarization calibrator (Category B with Note 3) with auxilary monitoring observations to transfer from primary. This calibration is needed to set the polarization vector angle 0.5*arctan(U/Q) and should be done in all cases. There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (Category C) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source (Category B) * Two scans of a source of known polarization (Category A or B with transfer) See Tables 1-4 below for Category A-D source catalogs. ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (<6cm) || 1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || 2 |- |} Notes: *1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0555+3948 || B0552+398 (3C138) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || 3 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || 3 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || 3 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || 4 |- |} Notes: *3. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *4. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || 5 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || 6 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || 7 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || 6 |- |} Notes: *5. Very bright and low polarization above 4GHz (but variable flux density). *6. Weak at high frequency, but stable flux and very low polarization. *7. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. effcc96f39cda26d1d1ee061d75a175255af7ead 191 190 2010-11-17T17:34:00Z Smyers 4 /* Polarization Calibrator Catalog and Selection */ wikitext text/x-wiki = Polarization Calibration = == Revised OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Detailed Guidelines == === Observing Recommendations === There are several strategies for deriving the Q/U angle calibration: * Observation of a primary polarization standard (Category A) * Observation of a secondary polarization calibrator (Category B with Note 3) with auxilary monitoring observations to transfer from primary. This calibration is needed to set the polarization vector angle 0.5*arctan(U/Q) and should be done in all cases. There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (Category C) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source (Category B) * Two scans of a source of known polarization (Category A or B with transfer) See Tables 1-4 below for Category A-D source catalogs. ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (<6cm) || 1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || 2 |- |} Notes: *1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0555+3948 || B0552+398 (3C138) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || 3 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || 3 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || 3 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || 4 |- |} Notes: *3. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *4. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || 5 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || 6 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || 7 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || 6 |- |} Notes: *5. Very bright and low polarization, but variable flux density. Approaches 1% polarized above 40GHz. *6. Weak at high frequency, but stable flux and very low polarization. *7. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. d9f54604664a9ba5149d456c745571be51acecea 192 191 2010-11-17T17:35:38Z Smyers 4 /* Polarization Calibrator Catalog and Selection */ wikitext text/x-wiki = Polarization Calibration = == Revised OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Detailed Guidelines == === Observing Recommendations === There are several strategies for deriving the Q/U angle calibration: * Observation of a primary polarization standard (Category A) * Observation of a secondary polarization calibrator (Category B with Note 3) with auxilary monitoring observations to transfer from primary. This calibration is needed to set the polarization vector angle 0.5*arctan(U/Q) and should be done in all cases. There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (Category C) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source (Category B) * Two scans of a source of known polarization (Category A or B with transfer) See Tables 1-4 below for Category A-D source catalogs. ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (>4GHz) || 1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || 2 |- |} Notes: *1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0555+3948 || B0552+398 (3C138) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || 3 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || 3 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || 3 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || 4 |- |} Notes: *3. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *4. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || 5 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || 6 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || 7 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || 6 |- |} Notes: *5. Very bright and low polarization, but variable flux density. Approaches 1% polarized above 40GHz. *6. Weak at high frequency, but stable flux and very low polarization. *7. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. a0a1e34e7656d40919b9d453dc2c75d96afed174 193 192 2010-11-17T17:47:51Z Smyers 4 /* Polarization Calibrator Catalog and Selection */ wikitext text/x-wiki = Polarization Calibration = == Revised OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Detailed Guidelines == === Observing Recommendations === There are several strategies for deriving the Q/U angle calibration: * Observation of a primary polarization standard (Category A) * Observation of a secondary polarization calibrator (Category B with Note 3) with auxilary monitoring observations to transfer from primary. This calibration is needed to set the polarization vector angle 0.5*arctan(U/Q) and should be done in all cases. There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (Category C) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source (Category B) * Two scans of a source of known polarization (Category A or B with transfer) See Tables 1-4 below for Category A-D source catalogs. ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (>4GHz) || 1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || 2 |- |} Notes: *1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0555+3948 || B0552+398 (3C138) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || 3 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || 3 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || 3 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || 4 |- |} Notes: *3. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *4. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || 5 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || 6 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || 7 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || 6 |- |} Notes: *5. Very bright and low polarization (<1%), but variable flux density. Approaches 1% polarized at 43GHz. *6. Weak at high frequency, but stable flux and very low polarization. *7. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. 54bd4ef64f12bfd94f2375e3b13cc8fbce9fd19e 14 21 195 2010-12-02T17:07:41Z Jmcmulli 2 Created page with "= Polarization Calibration = == Revised OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observati..." wikitext text/x-wiki = Polarization Calibration = == Revised OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Detailed Guidelines == === Observing Recommendations === There are several strategies for deriving the Q/U angle calibration: * Observation of a primary polarization standard (Category A) * Observation of a secondary polarization calibrator (Category B with Note 3) with auxilary monitoring observations to transfer from primary. This calibration is needed to set the polarization vector angle 0.5*arctan(U/Q) and should be done in all cases. There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (Category C) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source (Category B) * Two scans of a source of known polarization (Category A or B with transfer) See Tables 1-4 below for Category A-D source catalogs. ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (>4GHz) || 1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || 2 |- |} Notes: *1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0555+3948 || B0552+398 (3C138) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || 3 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || 3 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || 3 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || 4 |- |} Notes: *3. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *4. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || 5 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || 6 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || 7 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || 6 |- |} Notes: *5. Very bright and low polarization (<1%), but variable flux density. Approaches 1% polarized at 43GHz. *6. Weak at high frequency, but stable flux and very low polarization. *7. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === Perhaps just a pointer to the CASA guides page for the relevant section. 54bd4ef64f12bfd94f2375e3b13cc8fbce9fd19e 14 22 196 2010-12-02T17:23:55Z Jmcmulli 2 Created page with "= High Frequency Observing (K, Ka, Q) = == Current->Revised OSS Guidelines == * Sensitivity ** SEFD listed at 1400 Jy * EVLA Frequency Bands and Tunability ** In general, for a..." wikitext text/x-wiki = High Frequency Observing (K, Ka, Q) = == Current->Revised OSS Guidelines == * Sensitivity ** SEFD listed at 1400 Jy * EVLA Frequency Bands and Tunability ** In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: *** At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. *** At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. * Calibration and Flux Density Scale ** Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Note: Must update OSS Flux densities of Standard Calibrators * General Guidelines for Gain Calibration ** If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. * Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) ** For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Documentation. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Calibration Strategy ==== ==== Known RFI ==== ==== Calibrator Catalog and Selection ==== ==== Monitoring Observations ==== === Post-processing Guidelines === 7be4ff2bbc7beac6e7ee35c0803d41a955b29804 197 196 2010-12-02T17:24:13Z Jmcmulli 2 /* Current->Revised OSS Guidelines */ wikitext text/x-wiki = High Frequency Observing (K, Ka, Q) = == Current->Revised OSS Guidelines == * Sensitivity ** SEFD listed at 1400 Jy * EVLA Frequency Bands and Tunability ** In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: *** At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. *** At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. * Calibration and Flux Density Scale ** Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Note: Must update OSS Flux densities of Standard Calibrators * General Guidelines for Gain Calibration ** If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. * Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) ** For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Documentation. ** Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Calibration Strategy ==== ==== Known RFI ==== ==== Calibrator Catalog and Selection ==== ==== Monitoring Observations ==== === Post-processing Guidelines === 34babdcc973fd96f07bc994f038f1d88508e3199 198 197 2010-12-02T17:24:33Z Jmcmulli 2 /* Current->Revised OSS Guidelines */ wikitext text/x-wiki = High Frequency Observing (K, Ka, Q) = == Current->Revised OSS Guidelines == * Sensitivity ** SEFD listed at 1400 Jy * EVLA Frequency Bands and Tunability ** In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: *** At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. *** At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. * Calibration and Flux Density Scale ** Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Note: Must update OSS Flux densities of Standard Calibrators * General Guidelines for Gain Calibration ** If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. * Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) ** For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Documentation. ** Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Calibration Strategy ==== ==== Known RFI ==== ==== Calibrator Catalog and Selection ==== ==== Monitoring Observations ==== === Post-processing Guidelines === 5e724b20fa7684a017d85ae1854b7fbc7e341d1e 199 198 2010-12-02T18:34:27Z Jmcmulli 2 /* Detailed Guidelines */ wikitext text/x-wiki = High Frequency Observing (K, Ka, Q) = == Current->Revised OSS Guidelines == * Sensitivity ** SEFD listed at 1400 Jy * EVLA Frequency Bands and Tunability ** In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: *** At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. *** At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. * Calibration and Flux Density Scale ** Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Note: Must update OSS Flux densities of Standard Calibrators * General Guidelines for Gain Calibration ** If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. * Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) ** For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Documentation. ** Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== ==== Known RFI ==== ==== Calibrator Catalog and Selection ==== ==== Monitoring Observations ==== === Post-processing Guidelines === 18bdecd303f0a6a6cafbacb7a3fb0db7b073835f 6 24 205 2010-12-10T15:36:27Z Jmcmulli 2 Receiver Availability Forecast - 2010 Dec wikitext text/x-wiki Receiver Availability Forecast - 2010 Dec 2f2e29dd9d25fc33371734638e1040490bb2b6df 272 205 2011-06-29T15:14:37Z Jmcmulli 2 uploaded a new version of &quot;[[File:WideBandRcvrFrcstDec10.png]]&quot;: June 2011 update; includes 3-bit sampler schedule. wikitext text/x-wiki Receiver Availability Forecast - 2010 Dec 2f2e29dd9d25fc33371734638e1040490bb2b6df 10 2 211 200 2010-12-10T16:17:30Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] · [[Observational Status Summary - Projected]] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:SpectraLine| Spectral Line Observations]] · [[:Category:Polarimetry| Polarimetry Observations]] · [[:Category:Planetary| Planetary Observations]] · [[:Category:HighFrequency| High Frequency Observing (K, Ka, Q)]] · [[:Category:LowFrequency| Low Frequency Observing (L, S, C)]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] · [[Key to Calcodes]] |} 78bdb8a060ee12ae7abe506e50df7258ace2ddcb 228 211 2010-12-15T20:21:29Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] · [[Observational Status Summary]] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:SpectraLine| Spectral Line Observations]] · [[:Category:Polarimetry| Polarimetry Observations]] · [[:Category:Planetary| Planetary Observations]] · [[:Category:HighFrequency| High Frequency Observing (K, Ka, Q)]] · [[:Category:LowFrequency| Low Frequency Observing (L, S, C)]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] · [[Key to Calcodes]] |} de1e0e35226071b786199a2fd9f671fdf4f6a354 231 228 2010-12-15T20:26:17Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] · [[Observational Status Summary - Current]] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:SpectraLine| Spectral Line Observations]] · [[:Category:Polarimetry| Polarimetry Observations]] · [[:Category:Planetary| Planetary Observations]] · [[:Category:HighFrequency| High Frequency Observing (K, Ka, Q)]] · [[:Category:LowFrequency| Low Frequency Observing (L, S, C)]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] · [[Key to Calcodes]] |} 5806a7b69ee30a18a7687a591fd0db1986584a3e 232 231 2010-12-15T20:27:10Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:SpectraLine| Spectral Line Observations]] · [[:Category:Polarimetry| Polarimetry Observations]] · [[:Category:Planetary| Planetary Observations]] · [[:Category:HighFrequency| High Frequency Observing (K, Ka, Q)]] · [[:Category:LowFrequency| Low Frequency Observing (L, S, C)]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] · [[Key to Calcodes]] |} d1ca9a7e749d6b2fca2bc67c01d1372fc0920eff 249 232 2011-02-15T22:35:31Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:SpectraLine| Spectral Line Observations]] · [[:Category:Polarimetry| Polarimetry Observations]] · [[:Category:Planetary| Planetary Observations]] · [[:Category:HighFrequency| High Frequency Observing (K, Ka, Q)]] · [[:Category:LowFrequency| Low Frequency Observing (L, S, C)]] · [[:Category:PhasedArray| Phased Array Observing]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] · [[Key to Calcodes]] |} 66765c446bd42d02e2a5ae3b8ec7aa1a580d9354 255 249 2011-03-13T23:10:06Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [[:Category:Status|Observational Status Summary]] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:SpectraLine| Spectral Line Observations]] · [[:Category:Polarimetry| Polarimetry Observations]] · [[:Category:Planetary| Planetary Observations]] · [[:Category:HighFrequency| High Frequency Observing (K, Ka, Q)]] · [[:Category:LowFrequency| Low Frequency Observing (L, S, C)]] · [[:Category:PhasedArray| Phased Array Observing]] [[:Category:Pulsar| Pulsar Observing]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] · [[Key to Calcodes]] |} 8befb77371fd2f22e5de11cc5bde996c180bb7ba 6 15 214 119 2010-12-14T15:59:50Z Jmcmulli 2 uploaded a new version of &quot;[[File:SEFD.png]]&quot;: Adds color; Ku band points. wikitext text/x-wiki Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear. 73d0202b11dfa76074665912604c2f3fb1f47a52 215 214 2010-12-14T16:05:07Z Jmcmulli 2 uploaded a new version of &quot;[[File:SEFD.png]]&quot;: Adjust limits of lower frequency plot. wikitext text/x-wiki Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear. 73d0202b11dfa76074665912604c2f3fb1f47a52 273 215 2011-06-30T00:36:33Z Jmcmulli 2 uploaded a new version of &quot;[[File:SEFD.png]]&quot; wikitext text/x-wiki Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear. 73d0202b11dfa76074665912604c2f3fb1f47a52 0 1 225 151 2010-12-15T20:03:56Z Jmcmulli 2 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">Featured Article</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{Featured Article}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">News</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | * 02-Mar-2010: First science observations with WIDAR * 12-Apr-2010: First 27-antenna correlation * 25-Jun-2010: First fringes with 3-bit samplers (3 antennas) * 05-Jul-2010: First 2 GHz BW science observing * 09-Sep-2010: First 74 MHz observations with WIDAR * 13-Oct-2010: Time averaging enabled in correlator back-end {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:left;" | <div style="font-size:120%">Configuration & Proposal Dates</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| |- | style="padding-left:6px; padding-top:6px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| ! Trimester ! Observing Period ! Configuration ! Proposal Deadline |- | 2010 Mar 01 - 2010 Sep 13 | align="center" | D | 2009 Oct 1 |- | 2010 Sep 17 - 2010 Oct 04 | align="center" | DnC | 2009 Oct 1 |- | 2010 Oct 15 - 2011 Jan 18 | align="center" | C | 2010 Feb 1 |- | 2011 Jan 21 - 2011 Feb 07 | align="center" | CnB | 2010 Feb 1 |- | 2011 Feb 18 - 2011 May 09 | align="center" | B | 2010 Jun 1 |- | 2011 May 13 - 2011 May 31 | align="center" | BnA | 2010 Jun 1 |- | 2011 Jun 10 - 2011 Sep 12 | align='center' | A | 2010 Oct 1 |} <!-- TEXT 2 (END) --> |} |} 1dda164604a57accd6ca4d450320e45232e2282a 14 8 229 172 2010-12-15T20:24:49Z Jmcmulli 2 Replaced content with "= EVLA Observational Status = * [[Observational Status Summary]]: Appropriate for Proposal Calls * [[Observational Status Summary - Current]]: Appropriate for 'current' obse..." wikitext text/x-wiki = EVLA Observational Status = * [[Observational Status Summary]]: Appropriate for Proposal Calls * [[Observational Status Summary - Current]]: Appropriate for 'current' observations 2105f088592915e9db2203a1f0344c5ea12b857f 233 229 2010-12-15T20:27:39Z Jmcmulli 2 /* EVLA Observational Status */ wikitext text/x-wiki = EVLA Observational Status = * [[Observational Status Summary]]: Appropriate for Proposal Calls (updated 15 Dec 2010). * [[Observational Status Summary - Current]]: Appropriate for 'current' observations (updated 15 May 2010). 8813a465bcabdc2b0713bb6633e2902cde908790 238 233 2011-01-04T23:09:25Z Cchandle 6 wikitext text/x-wiki = EVLA Observational Status = * [[Observational Status Summary]]: Describes expected OSRO capabilities for the next D→A configuration cycle, September 2011 through January 2013, for use in preparing proposals for the February 1, 2011, proposal deadline, and following (updated 15 Dec 2010). * [[Observational Status Summary - Current]]: Describes current OSRO capabilities through the end of the A configuration, September 2011, for use by observers and rapid response science proposers (updated 15 May 2010). 6c5cb9de2a1113c902fe6fed89eb174fb7519b44 282 238 2011-06-30T16:44:01Z Jmcmulli 2 /* EVLA Observational Status */ wikitext text/x-wiki = EVLA Observational Status = * [[Observational Status Summary]]: Describes expected OSRO capabilities for the next D→A configuration cycle, September 2011 through January 2013, for use in preparing proposals for the August 1, 2011, proposal deadline, and following (updated 15 Dec 2010). * [[Observational Status Summary - Current]]: Describes current OSRO capabilities through the end of the A configuration, September 2011, for use by observers and rapid response science proposers (updated 15 May 2010). e11278d4741e6d172c67f6c90bd154898e25a0d9 283 282 2011-07-05T20:11:01Z Jmcmulli 2 /* EVLA Observational Status */ wikitext text/x-wiki = EVLA Observational Status = * [[Observational Status Summary]]: Describes expected OSRO capabilities for the next D→A configuration cycle, September 2011 through January 2013, for use in preparing proposals for the August 1, 2011, proposal deadline, and following (updated 29 Jun 2011). * [[Observational Status Summary - Current]]: Describes current OSRO capabilities through the end of the A configuration, September 2011, for use by observers and rapid response science proposers (updated 15 May 2010). be7680c970a07e390f2f22bcc29d125d77a0c16b 14 27 234 2010-12-15T21:57:20Z Jmcmulli 2 Created page with "= Low Frequency Observing (L, S, C) = == Current->Revised OSS Guidelines == * Sensitivity ** L band: 420 (RMS confusion level: 89) ** S band: 370 (RMS confusion level: 14) ** C..." wikitext text/x-wiki = Low Frequency Observing (L, S, C) = == Current->Revised OSS Guidelines == * Sensitivity ** L band: 420 (RMS confusion level: 89) ** S band: 370 (RMS confusion level: 14) ** C band: 310 (RMS confusion level: 2.3) * EVLA Frequency Bands and Tunability ** No special notes * Calibration and Flux Density Scale ** Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Note: Must update OSS Flux densities of Standard Calibrators; need S band fluxes. * General Guidelines for Gain Calibration ** If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== ==== Known RFI ==== ==== Calibrator Catalog and Selection ==== ==== Monitoring Observations ==== === Post-processing Guidelines === d08ef6e9466a4389d8165c580fecd96eead39f67 236 234 2010-12-16T00:16:42Z Jmcmulli 2 /* Current->Revised OSS Guidelines */ wikitext text/x-wiki = Low Frequency Observing (L, S, C) = == Current->Revised OSS Guidelines == * Sensitivity ** L band: 420 (RMS confusion level: 89) ** S band: 370 (RMS confusion level: 14) ** C band: 310 (RMS confusion level: 2.3) * EVLA Frequency Bands and Tunability ** No special notes * Calibration and Flux Density Scale ** LOOK AT THE NVSS ** Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Note: Must update OSS Flux densities of Standard Calibrators; need S band fluxes. * General Guidelines for Gain Calibration ** If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== ==== Known RFI ==== ==== Calibrator Catalog and Selection ==== ==== Monitoring Observations ==== === Post-processing Guidelines === a48ff3a2cadf0e6ebed4efc80d0a737dd2c883ee 14 21 247 195 2011-02-10T17:00:15Z Jmcmulli 2 /* Post-processing Guidelines */ wikitext text/x-wiki = Polarization Calibration = == Revised OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Detailed Guidelines == === Observing Recommendations === There are several strategies for deriving the Q/U angle calibration: * Observation of a primary polarization standard (Category A) * Observation of a secondary polarization calibrator (Category B with Note 3) with auxilary monitoring observations to transfer from primary. This calibration is needed to set the polarization vector angle 0.5*arctan(U/Q) and should be done in all cases. There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (Category C) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source (Category B) * Two scans of a source of known polarization (Category A or B with transfer) See Tables 1-4 below for Category A-D source catalogs. ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (>4GHz) || 1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || 2 |- |} Notes: *1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0555+3948 || B0552+398 (3C138) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || 3 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || 3 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || 3 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || 4 |- |} Notes: *3. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *4. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || 5 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || 6 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || 7 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || 6 |- |} Notes: *5. Very bright and low polarization (<1%), but variable flux density. Approaches 1% polarized at 43GHz. *6. Weak at high frequency, but stable flux and very low polarization. *7. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === For CASA reduction and analysis of polarization data, please see the following links: * [[http://casa.nrao.edu/docs/UserMan/UserMansu184.html#x213-2100004.4.5| Instrumental Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu185.html#x214-2110004.4.5.1| Heuristics and Strategies for Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu186.html#x215-2120004.4.5.2| A Note on Channelized Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu187.html#x216-2130004.4.5.3| A Polarization Calibration Example]] 654bcfb6dca3a9f4152970d2b593581e41b6961b 248 247 2011-02-10T17:01:51Z Jmcmulli 2 wikitext text/x-wiki = Polarization Calibration = == Revised OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Detailed Guidelines == === Observing Recommendations === There are several strategies for deriving the Q/U angle calibration: * Observation of a primary polarization standard (Category A) * Observation of a secondary polarization calibrator (Category B with Note 3) with auxilary monitoring observations to transfer from primary. This calibration is needed to set the polarization vector angle 0.5*arctan(U/Q) and should be done in all cases. There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (Category C) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source (Category B) * Two scans of a source of known polarization (Category A or B with transfer) See Tables 1-4 below for Category A-D source catalogs. ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Wide Band Considerations ==== * Performance over 1, 2 and 8 GHz band widths ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (>4GHz) || 1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || 2 |- |} Notes: *1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0555+3948 || B0552+398 (3C138) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || 3 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || 3 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || 3 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || 4 |- |} Notes: *3. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *4. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || 5 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || 6 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || 7 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || 6 |- |} Notes: *5. Very bright and low polarization (<1%), but variable flux density. Approaches 1% polarized at 43GHz. *6. Weak at high frequency, but stable flux and very low polarization. *7. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === For CASA reduction and analysis of polarization data, please see the following links: * [[http://casa.nrao.edu/docs/UserMan/UserMansu184.html#x213-2100004.4.5| Instrumental Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu185.html#x214-2110004.4.5.1| Heuristics and Strategies for Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu186.html#x215-2120004.4.5.2| A Note on Channelized Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu187.html#x216-2130004.4.5.3| A Polarization Calibration Example]] 8d2eac3d5f40f87121bc710e5672a7d80decd941 269 248 2011-05-10T22:26:01Z Smyers 4 /* Polarization Calibrator Catalog and Selection */ wikitext text/x-wiki = Polarization Calibration = == Revised OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Detailed Guidelines == === Observing Recommendations === There are several strategies for deriving the Q/U angle calibration: * Observation of a primary polarization standard (Category A) * Observation of a secondary polarization calibrator (Category B with Note 3) with auxilary monitoring observations to transfer from primary. This calibration is needed to set the polarization vector angle 0.5*arctan(U/Q) and should be done in all cases. There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (Category C) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source (Category B) * Two scans of a source of known polarization (Category A or B with transfer) See Tables 1-4 below for Category A-D source catalogs. ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Wide Band Considerations ==== * Performance over 1, 2 and 8 GHz band widths ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (>4GHz) || 1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || 2 |- |} Notes: *1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || 3 |- | J0555+3948 || B0552+398 || bright, flat spectrum, monitored, moderate variability || 3 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || 3 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || 3 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || 3 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || 4 |- |} Notes: *3. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *4. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || 5 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || 6 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || 7 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || 6 |- |} Notes: *5. Very bright and low polarization (<1%), but variable flux density. Approaches 1% polarized at 43GHz. *6. Weak at high frequency, but stable flux and very low polarization. *7. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === For CASA reduction and analysis of polarization data, please see the following links: * [[http://casa.nrao.edu/docs/UserMan/UserMansu184.html#x213-2100004.4.5| Instrumental Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu185.html#x214-2110004.4.5.1| Heuristics and Strategies for Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu186.html#x215-2120004.4.5.2| A Note on Channelized Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu187.html#x216-2130004.4.5.3| A Polarization Calibration Example]] 21bc6495fd73602d2f75179bed75a747b4f7415a 270 269 2011-05-10T22:42:02Z Smyers 4 /* Polarization Calibrator Catalog and Selection */ wikitext text/x-wiki = Polarization Calibration = == Revised OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Detailed Guidelines == === Observing Recommendations === There are several strategies for deriving the Q/U angle calibration: * Observation of a primary polarization standard (Category A) * Observation of a secondary polarization calibrator (Category B with Note 3) with auxilary monitoring observations to transfer from primary. This calibration is needed to set the polarization vector angle 0.5*arctan(U/Q) and should be done in all cases. There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (Category C) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source (Category B) * Two scans of a source of known polarization (Category A or B with transfer) See Tables 1-4 below for Category A-D source catalogs. ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Wide Band Considerations ==== * Performance over 1, 2 and 8 GHz band widths ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (>4GHz) || A1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || A2 |- |} Notes: *A1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *A2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || B1 |- | J0555+3948 || B0552+398 || bright, flat spectrum, monitored, moderate variability || B1,B2 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || B1 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || B1,B2 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || B1 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || B3 |- |} Notes: *B1. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *B2. Low polarization at low frequencies (L, sometimes S,C), do not use as angle calibrator. *B3. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || C1 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || C2 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || C3 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || C2 |- |} Notes: *C1. Very bright and low polarization (<1%), but variable flux density. Approaches 1% polarized at 43GHz. *C2. Weak at high frequency, but stable flux and very low polarization. *C3. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === For CASA reduction and analysis of polarization data, please see the following links: * [[http://casa.nrao.edu/docs/UserMan/UserMansu184.html#x213-2100004.4.5| Instrumental Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu185.html#x214-2110004.4.5.1| Heuristics and Strategies for Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu186.html#x215-2120004.4.5.2| A Note on Channelized Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu187.html#x216-2130004.4.5.3| A Polarization Calibration Example]] ee44ba0711aab7c8555dac8ad65b30117be12fb1 279 270 2011-06-30T16:16:29Z Jmcmulli 2 /* Revised OSS Guidelines */ wikitext text/x-wiki = Polarization Calibration = == Revised OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Detailed Guidelines == === Observing Recommendations === There are several strategies for deriving the Q/U angle calibration: * Observation of a primary polarization standard (Category A) * Observation of a secondary polarization calibrator (Category B with Note 3) with auxilary monitoring observations to transfer from primary. This calibration is needed to set the polarization vector angle 0.5*arctan(U/Q) and should be done in all cases. There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (Category C) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source (Category B) * Two scans of a source of known polarization (Category A or B with transfer) See Tables 1-4 below for Category A-D source catalogs. ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Wide Band Considerations ==== * Performance over 1, 2 and 8 GHz band widths ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (>4GHz) || A1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || A2 |- |} Notes: *A1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *A2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || B1 |- | J0555+3948 || B0552+398 || bright, flat spectrum, monitored, moderate variability || B1,B2 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || B1 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || B1,B2 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || B1 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || B3 |- |} Notes: *B1. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *B2. Low polarization at low frequencies (L, sometimes S,C), do not use as angle calibrator. *B3. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || C1 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || C2 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || C3 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || C2 |- |} Notes: *C1. Very bright and low polarization (<1%), but variable flux density. Approaches 1% polarized at 43GHz. *C2. Weak at high frequency, but stable flux and very low polarization. *C3. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === For CASA reduction and analysis of polarization data, please see the following links: * [[http://casa.nrao.edu/docs/UserMan/UserMansu184.html#x213-2100004.4.5| Instrumental Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu185.html#x214-2110004.4.5.1| Heuristics and Strategies for Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu186.html#x215-2120004.4.5.2| A Note on Channelized Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu187.html#x216-2130004.4.5.3| A Polarization Calibration Example]] 3560dec3f32c5400e2d5a6c40bd68b03fe27f527 14 28 250 2011-02-15T22:38:43Z Jmcmulli 2 Created page with "= Phased Array Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA Another mode of observing (commonly called phased array) will allow operation of the..." wikitext text/x-wiki = Phased Array Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it has not yet been commissioned. * VLBI Observing VLBI observations with the EVLA, such as phased array and single-antenna (e.g., Y1) modes, have not yet been commissioned and are not yet available to the OSRO program. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== ==== Known RFI ==== ==== Calibrator Catalog and Selection ==== ==== Monitoring Observations ==== === Post-processing Guidelines === 8fad7629768f4f56ce9611e794b2c38eb4613f4c 251 250 2011-02-15T22:39:18Z Jmcmulli 2 wikitext text/x-wiki = Phased Array Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it has not yet been commissioned. * VLBI Observing VLBI observations with the EVLA, such as phased array and single-antenna (e.g., Y1) modes, have not yet been commissioned and are not yet available to the OSRO program. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== === Post-processing Guidelines === ca788da9dbb22fb5a6621f94595b10d7c60980e5 14 29 252 2011-02-15T22:58:43Z Jmcmulli 2 Created page with "= Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Detai..." wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== ==== Monitoring Observations ==== === Post-processing Guidelines === 2e8cb866c386e6bf584125846255fa42b450fbda 253 252 2011-03-02T16:29:40Z Jott 8 wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== === Bandpass Setup === All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S$_{cal}$ observed for a time t$_{cal}$ and a science target with flux density S$_{obj}$ observed for a time t$_{obj}$, S$_{cal} \sqrt{t_cal}$ should be greater than S$_{obj} \sqrt{t_obj}$. How many times greater will be determined by one's science goals and the practicalities of the observations, but S$_{cal} \sqrt{t_cal}$ should be greater by at least a factor of two (??). The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of $\sim$0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few ($\sim$2--4) tenths of a percent over a period of several ($\sim$4--8) hours. [IS THIS TRUE FOR L BAND?] However, many antennas do show time-variable structure in the bandpass at low levels ($<$0.1\%). Dramatic jumps in the bandpass structure (of order a few percent) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any ``problem'' antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === e401efa316b9fa2f476a06049623514998cce83d 254 253 2011-03-02T16:31:41Z Jott 8 wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S$_{cal}$ observed for a time t$_{cal}$ and a science target with flux density S$_{obj}$ observed for a time t$_{obj}$, S$_{cal} \sqrt{t_cal}$ should be greater than S$_{obj} \sqrt{t_obj}$. How many times greater will be determined by one's science goals and the practicalities of the observations, but S$_{cal} \sqrt{t_cal}$ should be greater by at least a factor of two (??). The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of $\sim$0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few ($\sim$2--4) tenths of a percent over a period of several ($\sim$4--8) hours. [IS THIS TRUE FOR L BAND?] However, many antennas do show time-variable structure in the bandpass at low levels ($<$0.1\%). Dramatic jumps in the bandpass structure (of order a few percent) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any ``problem'' antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 0a015ae9bccbe1c21aa1084c9f23d89fac113789 257 254 2011-03-16T20:36:22Z Lchomiuk 9 /* Calibration Strategy */ wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S$_{cal}$ observed for a time t$_{cal}$ and a science target with flux density S$_{obj}$ observed for a time t$_{obj}$, S$_{cal} \sqrt{t_cal}$ should be greater than S$_{obj} \sqrt{t_obj}$. How many times greater will be determined by one's science goals and the practicalities of the observations, but S$_{cal} \sqrt{t_cal}$ should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of $\sim$0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few ($\sim$2--4) tenths of a percent over a period of several (~4--8) hours [L BAND?]. However, many antennas do show time-variable structure in the bandpass at low levels ($<$0.1\%). Dramatic jumps in the bandpass structure (of order a few percent) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any ``problem'' antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === bf2b333441e42a6d1bae6cc9cb32d8f616c3d266 258 257 2011-03-16T20:36:56Z Lchomiuk 9 /* Calibration Strategy */ wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S$_{cal}$ observed for a time t$_{cal}$ and a science target with flux density S$_{obj}$ observed for a time t$_{obj}$, S$_{cal} \sqrt{t_cal}$ should be greater than S$_{obj} \sqrt{t_obj}$. How many times greater will be determined by one's science goals and the practicalities of the observations, but S$_{cal} \sqrt{t_cal}$ should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of $\sim$0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few ($\sim$2--4) tenths of a percent over a period of several (~4--8) hours [L BAND?]. However, many antennas do show time-variable structure in the bandpass at low levels (<0.1%). Dramatic jumps in the bandpass structure (of order a few percent) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any ``problem'' antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. <math>Insert formula here</math> ==== Monitoring Observations ==== === Post-processing Guidelines === 50cbc6ef48e699ef99366042ebd53d9755b89263 259 258 2011-03-16T20:37:45Z Lchomiuk 9 /* Calibration Strategy */ wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S$_{cal}$ observed for a time t$_{cal}$ and a science target with flux density S$_{obj}$ observed for a time t$_{obj}$, S$_{cal} \sqrt{t_cal}$ should be greater than S$_{obj} \sqrt{t_obj}$. How many times greater will be determined by one's science goals and the practicalities of the observations, but S$_{cal} \sqrt{t_cal}$ should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) tenths of a percent over a period of several (~4--8) hours [L BAND?]. However, many antennas do show time-variable structure in the bandpass at low levels (<0.1%). Dramatic jumps in the bandpass structure (of order a few percent) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any ``problem'' antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === afb36403d6257e2765f5aa394036e6e8bd42333d 260 259 2011-03-16T20:38:12Z Lchomiuk 9 /* Calibration Strategy */ wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S$_{cal}$ observed for a time t$_{cal}$ and a science target with flux density S$_{obj}$ observed for a time t$_{obj}$, S$_{cal} \sqrt{t_cal}$ should be greater than S$_{obj} \sqrt{t_obj}$. How many times greater will be determined by one's science goals and the practicalities of the observations, but S$_{cal} \sqrt{t_cal}$ should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) tenths of a percent over a period of several (~4--8) hours [L BAND?]. However, many antennas do show time-variable structure in the bandpass at low levels (<0.1%). Dramatic jumps in the bandpass structure (of order a few percent) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === f31a16c18b27b7632a54de17355efa9ef577c88f 261 260 2011-03-16T20:41:14Z Lchomiuk 9 /* Calibration Strategy */ wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> $_{cal}$ observed for a time t$_{cal}$ and a science target with flux density S$_{obj}$ observed for a time t$_{obj}$, S$_{cal} \sqrt{t_cal}$ should be greater than S$_{obj} \sqrt{t_obj}$. How many times greater will be determined by one's science goals and the practicalities of the observations, but S$_{cal} \sqrt{t_cal}$ should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) tenths of a percent over a period of several (~4--8) hours [L BAND?]. However, many antennas do show time-variable structure in the bandpass at low levels (<0.1%). Dramatic jumps in the bandpass structure (of order a few percent) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 1f22f5351bd4bccf5ddd71aa1b17f1a35c306473 262 261 2011-03-16T20:41:59Z Lchomiuk 9 /* Calibration Strategy */ wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, S$_{cal} \sqrt{t_cal}$ should be greater than S$_{obj} \sqrt{t_obj}$. How many times greater will be determined by one's science goals and the practicalities of the observations, but S$_{cal} \sqrt{t_cal}$ should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) tenths of a percent over a period of several (~4--8) hours [L BAND?]. However, many antennas do show time-variable structure in the bandpass at low levels (<0.1%). Dramatic jumps in the bandpass structure (of order a few percent) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 907b291eca7e00f0ded45bce6a3b51e7b49a84ba 263 262 2011-03-16T20:43:16Z Lchomiuk 9 /* Calibration Strategy */ wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, S<sub>cal</sub> \sqrt{t<sub>cal</sub>} should be greater than S<sub>obj</sub> \sqrt{t<sub>obj</sub>}$. How many times greater will be determined by one's science goals and the practicalities of the observations, but S<sub>cal</sub> \sqrt{t<sub>cal</sub>}$ should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) tenths of a percent over a period of several (~4--8) hours [L BAND?]. However, many antennas do show time-variable structure in the bandpass at low levels (<0.1%). Dramatic jumps in the bandpass structure (of order a few percent) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 0c6433b832bd466507486b056edb09651f5112d1 264 263 2011-03-16T20:44:23Z Lchomiuk 9 /* Calibration Strategy */ wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, S<sub>cal</sub> \sqrt{t<sub>cal</sub>} should be greater than <math>S<sub>obj</sub> \sqrt{t<sub>obj</sub>}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but S<sub>cal</sub> \sqrt{t<sub>cal</sub>}$ should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) tenths of a percent over a period of several (~4--8) hours [L BAND?]. However, many antennas do show time-variable structure in the bandpass at low levels (<0.1%). Dramatic jumps in the bandpass structure (of order a few percent) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 06a18c5872295ff988170bc71804c75d65aa8c64 265 264 2011-03-16T20:46:25Z Lchomiuk 9 /* Calibration Strategy */ wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math>S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but S<sub>cal</sub> \sqrt{t<sub>cal</sub>}$ should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) tenths of a percent over a period of several (~4--8) hours [L BAND?]. However, many antennas do show time-variable structure in the bandpass at low levels (<0.1%). Dramatic jumps in the bandpass structure (of order a few percent) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 0dad9a62db32c392cc826c3938a026eb7610c7ae 266 265 2011-03-16T20:47:19Z Lchomiuk 9 /* Calibration Strategy */ wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) tenths of a percent over a period of several (~4--8) hours [L BAND?]. However, many antennas do show time-variable structure in the bandpass at low levels (<0.1%). Dramatic jumps in the bandpass structure (of order a few percent) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 02895b418e0578e7810f10ddb23505d6735e4bce 267 266 2011-03-16T21:32:21Z Lchomiuk 9 /* Calibration Strategy */ wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. However, many antennas do show time-variable structure in the bandpass at low levels (less than one part in a thousand). Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 22335a43a5e7f5c456e5ae6b9f2717a5adb887a6 268 267 2011-03-16T21:39:33Z Lchomiuk 9 /* Calibration Strategy */ wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 49639c89d3c94f52c6f9dd3ca8f396d07746360d 14 30 256 2011-03-13T23:11:08Z Jmcmulli 2 Created page with "= Pulsar Observing = == Current->Revised OSS Guidelines == * Note: No current guidelines in the OSS; we need to add basic information on setup and capabilities to guide proposa..." wikitext text/x-wiki = Pulsar Observing = == Current->Revised OSS Guidelines == * Note: No current guidelines in the OSS; we need to add basic information on setup and capabilities to guide proposal writing. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== === Post-processing Guidelines === de8b6b2511b48be47ff512084c49600535dba639 6 31 289 2011-07-13T22:13:55Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 32 290 2011-07-13T22:14:20Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 291 290 2011-07-13T22:14:44Z Jmcmulli 2 uploaded a new version of &quot;[[File:APIwind February.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 33 292 2011-07-13T22:14:55Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 34 293 2011-07-13T22:15:07Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 35 294 2011-07-13T22:15:24Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 36 295 2011-07-13T22:16:12Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 37 296 2011-07-13T22:16:29Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 38 297 2011-07-13T22:16:43Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 39 298 2011-07-13T22:17:00Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 40 299 2011-07-13T22:17:19Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 41 301 2011-07-13T22:33:16Z Jmcmulli 2 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 14 22 304 199 2011-08-18T15:30:44Z Cbrogan 11 /* Detailed Guidelines */ wikitext text/x-wiki = High Frequency Observing (K, Ka, Q) = == Current->Revised OSS Guidelines == * Sensitivity ** SEFD listed at 1400 Jy * EVLA Frequency Bands and Tunability ** In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: *** At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. *** At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. * Calibration and Flux Density Scale ** Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Note: Must update OSS Flux densities of Standard Calibrators * General Guidelines for Gain Calibration ** If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. * Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) ** For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Documentation. ** Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Detailed Guidelines == [[Draft_High_Freq_Details]] === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== ==== Known RFI ==== ==== Calibrator Catalog and Selection ==== ==== Monitoring Observations ==== === Post-processing Guidelines === 17bbcf305dea81f03f7bf616f8ad8916b432d6da 417 304 2011-08-25T15:33:48Z Cbrogan 11 wikitext text/x-wiki = High Frequency Observing (K, Ka, Q) = == Current->Revised OSS Guidelines == * Sensitivity ** SEFD listed at 1400 Jy * EVLA Frequency Bands and Tunability ** In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: *** At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. *** At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. * Calibration and Flux Density Scale ** Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Note: Must update OSS Flux densities of Standard Calibrators * General Guidelines for Gain Calibration ** If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. * Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) ** For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Documentation. ** Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Detailed Guidelines == [[High Frequency Observing][Draft by cbrogan]] === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== ==== Known RFI ==== ==== Calibrator Catalog and Selection ==== ==== Monitoring Observations ==== === Post-processing Guidelines === e75fde1c8b90f352fbeffde61dc3ba0fcb7b78eb 418 417 2011-08-25T15:35:26Z Cbrogan 11 /* Detailed Guidelines */ wikitext text/x-wiki = High Frequency Observing (K, Ka, Q) = == Current->Revised OSS Guidelines == * Sensitivity ** SEFD listed at 1400 Jy * EVLA Frequency Bands and Tunability ** In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: *** At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. *** At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. * Calibration and Flux Density Scale ** Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Note: Must update OSS Flux densities of Standard Calibrators * General Guidelines for Gain Calibration ** If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. * Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) ** For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Documentation. ** Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Detailed Guidelines == [High Frequency Observing] Draft by C. Brogan === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== ==== Known RFI ==== ==== Calibrator Catalog and Selection ==== ==== Monitoring Observations ==== === Post-processing Guidelines === ce76d471c181f7e0512bf3f4f4cc4b0c3c702344 419 418 2011-08-25T15:35:50Z Cbrogan 11 /* Detailed Guidelines */ wikitext text/x-wiki = High Frequency Observing (K, Ka, Q) = == Current->Revised OSS Guidelines == * Sensitivity ** SEFD listed at 1400 Jy * EVLA Frequency Bands and Tunability ** In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: *** At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. *** At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. * Calibration and Flux Density Scale ** Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Note: Must update OSS Flux densities of Standard Calibrators * General Guidelines for Gain Calibration ** If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. * Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) ** For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Documentation. ** Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Detailed Guidelines == [[High Frequency Observing]] Draft by C. Brogan === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== ==== Known RFI ==== ==== Calibrator Catalog and Selection ==== ==== Monitoring Observations ==== === Post-processing Guidelines === 0170231f841d5a861863dbb3abd597f3e0ec4ab2 474 419 2011-08-31T23:32:19Z Thunter 10 /* Scheduling */ wikitext text/x-wiki = High Frequency Observing (K, Ka, Q) = == Current->Revised OSS Guidelines == * Sensitivity ** SEFD listed at 1400 Jy * EVLA Frequency Bands and Tunability ** In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: *** At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. *** At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. * Calibration and Flux Density Scale ** Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Note: Must update OSS Flux densities of Standard Calibrators * General Guidelines for Gain Calibration ** If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. * Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) ** For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Documentation. ** Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/. == Detailed Guidelines == [[High Frequency Observing]] Draft by C. Brogan === Observing Preparation Recommendations === ==== Scheduling ==== The scheduling of high frequency projects is dependent on the weather conditions. The two quantities currently used to determine if atmospheric conditions are suitable are the wind speed and the API rms phase. The fraction of time that both quantities are satisfactory depends on the time of year and time of day. A good summary of the monthly conditions at the EVLA as a function of LST is at this page: [[Monthly_Conditions_at_EVLA]]. ==== Calibration Strategy ==== ==== Known RFI ==== ==== Calibrator Catalog and Selection ==== ==== Monitoring Observations ==== === Post-processing Guidelines === 56f40eb3b1af9f3c3b087b2bb64f7fb3a7a4c99d 0 43 310 2011-08-18T19:44:27Z Thunter 10 Created page with "{| border="1" |+ RMS phase (in degrees) |- ! Month | rowspan="2" | 7mm | colspan="3" | 1.3cm | colspan="3" |- ! !! tcycle=2 min !! tcycle=5 min !! tcycle=10 min !!..." wikitext text/x-wiki {| border="1" |+ RMS phase (in degrees) |- ! Month | rowspan="2" | 7mm | colspan="3" | 1.3cm | colspan="3" |- ! !! tcycle=2 min !! tcycle=5 min !! tcycle=10 min !! 1.3 cm tcycle=4min !! tcycle=10 min !! tcycle=20 min |- | Day .. Night || Day .. Night || Day .. Night || Day .. Night || Day ..Night || Day .. Night |- ! Jan | 25.5 .. 18.8 || 47.1 .. 34.6 || 74.7 .. 54.9 || 25.5 .. 18.8 || 47.1 ..34.6 || 74.7 .. 54.9 |- ! Feb | 25.0 .. 15.1 || 46.1 .. 27.8 || 73.2 .. 44.2 || 25.0 .. 15.1 || 46.1 ..27.8 || 73.2 .. 44.2 |- ! Mar | 30.2 .. 20.9 || 55.7 .. 38.4 || 88.4 .. 61.0 || 30.2 .. 20.9 || 55.7 ..38.4 || 88.4 .. 61.0 |- ! Apr | 45.4 .. 24.5 || 83.5 .. 45.1 || 132.6 .. 71.6 || 45.4 .. 24.5 || 83.5 ..45.1 || 132.6 .. 71.6 |- ! May | 38.1 .. 21.9 || 70.1 .. 40.3 || 111.3 .. 64.0 || 38.1 .. 21.9 || 70.1 ..40.3 || 111.3 .. 64.0 |- ! Jun | 39.1 .. 20.3 || 72.0 .. 37.5 || 114.3 .. 59.5 || 39.1 .. 20.3 || 72.0 ..37.5 || 114.3 .. 59.5 |- ! Jul | 41.7 .. 27.6 || 76.8 .. 59.0 || 122.0 .. 80.8 || 39.1 .. 20.3 || 72.0 ..37.5 || 114.3 .. 59.5 |- ! Aug | 51.1 .. 29.7 || 94.1 .. 54.7 || 149.4 .. 86.9 || 51.1 .. 29.7 || 94.1 ..54.7 || 149.4 .. 86.9 |- ! Sep | 62.0 .. 34.9 || 114.3 .. 64.3 || 181.4 .. 102.1 || 62.0 .. 34.9 || 114.3 ..64.3 || 181.4 .. 102.1 |- ! Oct | 51.6 .. 30.8 || 95.1 .. 56.7 || 150.9 .. 89.9 || 51.6 .. 30.8 || 95.1 ..56.7 || 150.9 .. 89.9 |- ! Nov | 32.3 .. 29.7 || 59.5 .. 54.7 || 94.5 .. 86.9 || 32.3 .. 29.7 || 59.5 ..54.7 || 94.5 .. 86.9 |- ! Dec | 22.4 .. 15.6 || 41.3 .. 28.8 || 65.5 .. 45.7 || 22.4 .. 15.6 || 41.3 ..28.8 || 65.5 .. 45.7 |} 3f82a96befdbb23accd2a06f16344eccdb082ba4 312 310 2011-08-18T19:56:24Z Thunter 10 wikitext text/x-wiki {| border="1" class="wikitable" |+ <b>RMS phase (in degrees)</b> |- ! rowspan="3" align="center" | Month ! colspan="3" align="center" | 7mm ! colspan="3" align="center" | 1.3cm |- ! tcycle=2 min !! tcycle=5 min !! tcycle=10 min !! tcycle=4min !! tcycle=10 min !! tcycle=20 min |- ! Day .. Night !! Day .. Night !! Day .. Night !! Day .. Night !! Day ..Night !! Day .. Night |- ! Jan | 25.5 .. 18.8 || 47.1 .. 34.6 || 74.7 .. 54.9 || 25.5 .. 18.8 || 47.1 ..34.6 || 74.7 .. 54.9 |- ! Feb | 25.0 .. 15.1 || 46.1 .. 27.8 || 73.2 .. 44.2 || 25.0 .. 15.1 || 46.1 ..27.8 || 73.2 .. 44.2 |- ! Mar | 30.2 .. 20.9 || 55.7 .. 38.4 || 88.4 .. 61.0 || 30.2 .. 20.9 || 55.7 ..38.4 || 88.4 .. 61.0 |- ! Apr | 45.4 .. 24.5 || 83.5 .. 45.1 || 132.6 .. 71.6 || 45.4 .. 24.5 || 83.5 ..45.1 || 132.6 .. 71.6 |- ! May | 38.1 .. 21.9 || 70.1 .. 40.3 || 111.3 .. 64.0 || 38.1 .. 21.9 || 70.1 ..40.3 || 111.3 .. 64.0 |- ! Jun | 39.1 .. 20.3 || 72.0 .. 37.5 || 114.3 .. 59.5 || 39.1 .. 20.3 || 72.0 ..37.5 || 114.3 .. 59.5 |- ! Jul | 41.7 .. 27.6 || 76.8 .. 59.0 || 122.0 .. 80.8 || 39.1 .. 20.3 || 72.0 ..37.5 || 114.3 .. 59.5 |- ! Aug | 51.1 .. 29.7 || 94.1 .. 54.7 || 149.4 .. 86.9 || 51.1 .. 29.7 || 94.1 ..54.7 || 149.4 .. 86.9 |- ! Sep | 62.0 .. 34.9 || 114.3 .. 64.3 || 181.4 .. 102.1 || 62.0 .. 34.9 || 114.3 ..64.3 || 181.4 .. 102.1 |- ! Oct | 51.6 .. 30.8 || 95.1 .. 56.7 || 150.9 .. 89.9 || 51.6 .. 30.8 || 95.1 ..56.7 || 150.9 .. 89.9 |- ! Nov | 32.3 .. 29.7 || 59.5 .. 54.7 || 94.5 .. 86.9 || 32.3 .. 29.7 || 59.5 ..54.7 || 94.5 .. 86.9 |- ! Dec | 22.4 .. 15.6 || 41.3 .. 28.8 || 65.5 .. 45.7 || 22.4 .. 15.6 || 41.3 ..28.8 || 65.5 .. 45.7 |} 5b155db8804e2c7bc1427fa1a9b7ac9fcb62d2fb 313 312 2011-08-18T19:59:50Z Thunter 10 wikitext text/x-wiki {| border="1" class="wikitable" |+ <b>RMS phase (in degrees)</b> |- ! rowspan="1" align="center" | Wavelength: ! colspan="3" align="center" | 7 mm ! colspan="3" align="center" | 1.3 cm |- ! Cycle time: !! 2 min !! 5 min !! 10 min !! 4min !! 10 min !! 20 min |- ! Month !! Day .. Night !! Day .. Night !! Day .. Night !! Day .. Night !! Day ..Night !! Day .. Night |- ! Jan | 25.5 .. 18.8 || 47.1 .. 34.6 || 74.7 .. 54.9 || 25.5 .. 18.8 || 47.1 ..34.6 || 74.7 .. 54.9 |- ! Feb | 25.0 .. 15.1 || 46.1 .. 27.8 || 73.2 .. 44.2 || 25.0 .. 15.1 || 46.1 ..27.8 || 73.2 .. 44.2 |- ! Mar | 30.2 .. 20.9 || 55.7 .. 38.4 || 88.4 .. 61.0 || 30.2 .. 20.9 || 55.7 ..38.4 || 88.4 .. 61.0 |- ! Apr | 45.4 .. 24.5 || 83.5 .. 45.1 || 132.6 .. 71.6 || 45.4 .. 24.5 || 83.5 ..45.1 || 132.6 .. 71.6 |- ! May | 38.1 .. 21.9 || 70.1 .. 40.3 || 111.3 .. 64.0 || 38.1 .. 21.9 || 70.1 ..40.3 || 111.3 .. 64.0 |- ! Jun | 39.1 .. 20.3 || 72.0 .. 37.5 || 114.3 .. 59.5 || 39.1 .. 20.3 || 72.0 ..37.5 || 114.3 .. 59.5 |- ! Jul | 41.7 .. 27.6 || 76.8 .. 59.0 || 122.0 .. 80.8 || 39.1 .. 20.3 || 72.0 ..37.5 || 114.3 .. 59.5 |- ! Aug | 51.1 .. 29.7 || 94.1 .. 54.7 || 149.4 .. 86.9 || 51.1 .. 29.7 || 94.1 ..54.7 || 149.4 .. 86.9 |- ! Sep | 62.0 .. 34.9 || 114.3 .. 64.3 || 181.4 .. 102.1 || 62.0 .. 34.9 || 114.3 ..64.3 || 181.4 .. 102.1 |- ! Oct | 51.6 .. 30.8 || 95.1 .. 56.7 || 150.9 .. 89.9 || 51.6 .. 30.8 || 95.1 ..56.7 || 150.9 .. 89.9 |- ! Nov | 32.3 .. 29.7 || 59.5 .. 54.7 || 94.5 .. 86.9 || 32.3 .. 29.7 || 59.5 ..54.7 || 94.5 .. 86.9 |- ! Dec | 22.4 .. 15.6 || 41.3 .. 28.8 || 65.5 .. 45.7 || 22.4 .. 15.6 || 41.3 ..28.8 || 65.5 .. 45.7 |} 824bd9d8dc78595009ad3d72b4ae3d040aa3d556 314 313 2011-08-18T20:03:02Z Thunter 10 wikitext text/x-wiki {| border="1" class="wikitable" style="text-align: center;" |+ <b>RMS phase (in degrees)</b> |- ! rowspan="1" | Wavelength: ! colspan="3" | 7 mm ! colspan="3" | 1.3 cm |- ! Cycle time: !! 2 min !! 5 min !! 10 min !! 4min !! 10 min !! 20 min |- ! Month !! Day .. Night !! Day .. Night !! Day .. Night !! Day .. Night !! Day ..Night !! Day .. Night |- ! Jan | 25.5 .. 18.8 || 47.1 .. 34.6 || 74.7 .. 54.9 || 25.5 .. 18.8 || 47.1 ..34.6 || 74.7 .. 54.9 |- ! Feb | 25.0 .. 15.1 || 46.1 .. 27.8 || 73.2 .. 44.2 || 25.0 .. 15.1 || 46.1 ..27.8 || 73.2 .. 44.2 |- ! Mar | 30.2 .. 20.9 || 55.7 .. 38.4 || 88.4 .. 61.0 || 30.2 .. 20.9 || 55.7 ..38.4 || 88.4 .. 61.0 |- ! Apr | 45.4 .. 24.5 || 83.5 .. 45.1 || 132.6 .. 71.6 || 45.4 .. 24.5 || 83.5 ..45.1 || 132.6 .. 71.6 |- ! May | 38.1 .. 21.9 || 70.1 .. 40.3 || 111.3 .. 64.0 || 38.1 .. 21.9 || 70.1 ..40.3 || 111.3 .. 64.0 |- ! Jun | 39.1 .. 20.3 || 72.0 .. 37.5 || 114.3 .. 59.5 || 39.1 .. 20.3 || 72.0 ..37.5 || 114.3 .. 59.5 |- ! Jul | 41.7 .. 27.6 || 76.8 .. 59.0 || 122.0 .. 80.8 || 39.1 .. 20.3 || 72.0 ..37.5 || 114.3 .. 59.5 |- ! Aug | 51.1 .. 29.7 || 94.1 .. 54.7 || 149.4 .. 86.9 || 51.1 .. 29.7 || 94.1 ..54.7 || 149.4 .. 86.9 |- ! Sep | 62.0 .. 34.9 || 114.3 .. 64.3 || 181.4 .. 102.1 || 62.0 .. 34.9 || 114.3 ..64.3 || 181.4 .. 102.1 |- ! Oct | 51.6 .. 30.8 || 95.1 .. 56.7 || 150.9 .. 89.9 || 51.6 .. 30.8 || 95.1 ..56.7 || 150.9 .. 89.9 |- ! Nov | 32.3 .. 29.7 || 59.5 .. 54.7 || 94.5 .. 86.9 || 32.3 .. 29.7 || 59.5 ..54.7 || 94.5 .. 86.9 |- ! Dec | 22.4 .. 15.6 || 41.3 .. 28.8 || 65.5 .. 45.7 || 22.4 .. 15.6 || 41.3 ..28.8 || 65.5 .. 45.7 |} b87ae5fd925bc3222efca2a7ab7a73e6e330929a 6 44 380 2011-08-24T20:01:08Z Cbrogan 11 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 45 408 2011-08-24T22:13:39Z Cbrogan 11 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 46 413 2011-08-24T22:22:35Z Cbrogan 11 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 14 29 421 268 2011-08-26T06:25:34Z Jott 8 wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === [[JuergensSandbox]] 2b87fded4ccb2302c6c4e0b554998f407c18197a 423 421 2011-08-26T06:30:54Z Jott 8 wikitext text/x-wiki = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 49639c89d3c94f52c6f9dd3ca8f396d07746360d 0 48 422 2011-08-26T06:30:17Z Jott 8 Created page with "spectral line setup" wikitext text/x-wiki spectral line setup e7a038880bf700c59ea8c1015722b385542ec5cd 0 49 472 2011-08-31T23:26:59Z Thunter 10 Created page with "August 2011 [[Image:rmsphase_August2011.png|center|600px]]''API rms phase for August 2011''" wikitext text/x-wiki August 2011 [[Image:rmsphase_August2011.png|center|600px]]''API rms phase for August 2011'' d3f48271b9fa542035bbdaec28aaf568c8fc9add 477 472 2011-09-01T15:16:23Z Thunter 10 wikitext text/x-wiki The magenta line is the median computed per LST hour. [[Image:rmsphase_January2011.png|center|600px]]''January'' [[Image:rmsphase_February2011.png|center|600px]]'' '' [[Image:rmsphase_March2011.png|center|600px]]'' '' [[Image:rmsphase_April2011.png|center|600px]]'' '' [[Image:rmsphase_May2011.png|center|600px]]'' '' [[Image:rmsphase_June2011.png|center|600px]]'' '' [[Image:rmsphase_July2011.png|center|600px]]'' '' [[Image:rmsphase_August2011.png|center|600px]]'' '' [[Image:rmsphase_September2010.png|center|600px]]'' '' [[Image:rmsphase_October2010.png|center|600px]]'' '' [[Image:rmsphase_November2010.png|center|600px]]'' '' [[Image:rmsphase_December2010.png|center|600px]]'''' 034ee7ef02371e8fcb38657f362de848b6e0d7e4 481 477 2011-09-01T15:19:35Z Thunter 10 wikitext text/x-wiki The magenta line is the median computed per LST hour. [[Image:rmsphase_January2011.png|center|300px]]''January'' [[Image:rmsphase_February2011.png|center|300px]]''February''<br> [[Image:rmsphase_March2011.png|center|600px]]'' '' [[Image:rmsphase_April2011.png|center|600px]]'' '' [[Image:rmsphase_May2011.png|center|600px]]'' '' [[Image:rmsphase_June2011.png|center|600px]]'' '' [[Image:rmsphase_July2011.png|center|600px]]'' '' [[Image:rmsphase_August2011.png|center|600px]]'' '' [[Image:rmsphase_September2010.png|center|600px]]'' '' [[Image:rmsphase_October2010.png|center|600px]]'' '' [[Image:rmsphase_November2010.png|center|600px]]'' '' [[Image:rmsphase_December2010.png|center|600px]]'''' 3f064a3d00e1d0683e2bdd441ff56c8bfacaa61c 488 481 2011-09-01T15:41:20Z Thunter 10 wikitext text/x-wiki The magenta line is the median computed per LST hour. Click a plot to enlarge it. {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March ! [[Image:rmsphase_January2011.png|center|200px]]'' '' !! [[Image:rmsphase_February2011.png|center|200px]]'''' !! [[Image:rmsphase_March2011.png|center|200px]]'' '' |- ! April !! May !! June ! [[Image:rmsphase_April2011.png|center|200px]]'' '' || [[Image:rmsphase_May2011.png|center|200px]]'' '' || [[Image:rmsphase_June2011.png|center|200px]]'' '' |- ! July !! August !! September ! [[Image:rmsphase_July2011.png|center|600px]]'' '' || [[Image:rmsphase_August2011.png|center|200px]]'' '' || [[Image:rmsphase_September2010.png|center|200px]]'' '' |- ! October !! November !! December ! [[Image:rmsphase_October2010.png|center|200px]]'' '' || [[Image:rmsphase_November2010.png|center|200px]]'' '' || [[Image:rmsphase_December2010.png|center|200px]]'''' |} 0268e2149c5824976b8da2cf52c67da942efd9d4 489 488 2011-09-01T15:44:03Z Thunter 10 wikitext text/x-wiki The magenta line is the median computed per LST hour. Click a plot to enlarge it. {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rmsphase_January2011.png|center|200px]]'' '' !! [[Image:rmsphase_February2011.png|center|200px]]'''' !! [[Image:rmsphase_March2011.png|center|200px]]'' '' |- ! April !! May !! June |- | [[Image:rmsphase_April2011.png|center|200px]]'' '' || [[Image:rmsphase_May2011.png|center|200px]]'' '' || [[Image:rmsphase_June2011.png|center|200px]]'' '' |- ! July !! August !! September |- | [[Image:rmsphase_July2011.png|center|600px]]'' '' || [[Image:rmsphase_August2011.png|center|200px]]'' '' || [[Image:rmsphase_September2010.png|center|200px]]'' '' |- ! October !! November !! December |- | [[Image:rmsphase_October2010.png|center|200px]]'' '' || [[Image:rmsphase_November2010.png|center|200px]]'' '' || [[Image:rmsphase_December2010.png|center|200px]]'''' |} f56db510994896faa07fbc6395a30386a682c3bd 490 489 2011-09-01T15:51:23Z Thunter 10 wikitext text/x-wiki The magenta line is the median computed per LST hour. Click a plot to enlarge it. {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rmsphase_January2011.png|center|250px]]'' '' | [[Image:rmsphase_February2011.png|center|250px]]'' '' | [[Image:rmsphase_March2011.png|center|250px]]'' '' |- ! April !! May !! June |- | [[Image:rmsphase_April2011.png|center|250px]]'' '' | [[Image:rmsphase_May2011.png|center|250px]]'' '' | [[Image:rmsphase_June2011.png|center|250px]]'' '' |- ! July !! August !! September |- | [[Image:rmsphase_July2011.png|center|250px]]'' '' | [[Image:rmsphase_August2011.png|center|250px]]'' '' | [[Image:rmsphase_September2010.png|center|250px]]'' '' |- ! October !! November !! December |- | [[Image:rmsphase_October2010.png|center|250px]]'' '' | [[Image:rmsphase_November2010.png|center|250px]]'' '' | [[Image:rmsphase_December2010.png|center|250px]]'' '' |} 361636838828e3408ff55f4a3b22aee56f7b6d0e 6 50 473 2011-08-31T23:27:31Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 0 51 475 2011-08-31T23:33:08Z Thunter 10 Created page with "August 2011 [[Image:rx_August2011.png|center|600px]]''Fraction of time suitable for observing in August.''" wikitext text/x-wiki August 2011 [[Image:rx_August2011.png|center|600px]]''Fraction of time suitable for observing in August.'' 45b8a891ef24e0054eb6a0fc6d559c95b10f93c7 497 475 2011-09-01T15:59:42Z Thunter 10 wikitext text/x-wiki Fraction of time suitable for observing in each EVLA band. Click a plot to enlarge it. {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rx_January2011.png|center|250px]]'' '' | [[Image:rx_February2011.png|center|250px]]'' '' | [[Image:rx_March2011.png|center|250px]]'' '' |- ! April !! May !! June |- | [[Image:rx_April2011.png|center|250px]]'' '' | [[Image:rx_May2011.png|center|250px]]'' '' | [[Image:rx_June2011.png|center|250px]]'' '' |- ! July !! August !! September |- | [[Image:rx_July2011.png|center|250px]]'' '' | [[Image:rx_August2011.png|center|250px]]'' '' | [[Image:rx_September2010.png|center|250px]]'' '' |- ! October !! November !! December |- | [[Image:rx_October2010.png|center|250px]]'' '' | [[Image:rx_November2010.png|center|250px]]'' '' | [[Image:rx_December2010.png|center|250px]]'' '' |} a825befc2694f07d58ba605f249a6e65bbf328f7 6 52 476 2011-08-31T23:33:20Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 53 478 2011-09-01T15:17:00Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 479 478 2011-09-01T15:17:28Z Thunter 10 uploaded a new version of &quot;[[File:Rmsphase January2011.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 494 479 2011-09-01T15:55:22Z Thunter 10 uploaded a new version of &quot;[[File:Rmsphase January2011.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 54 480 2011-09-01T15:18:04Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 495 480 2011-09-01T15:56:10Z Thunter 10 uploaded a new version of &quot;[[File:Rmsphase February2011.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 55 482 2011-09-01T15:20:13Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 496 482 2011-09-01T15:56:37Z Thunter 10 uploaded a new version of &quot;[[File:Rmsphase March2011.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 56 483 2011-09-01T15:20:54Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 57 484 2011-09-01T15:22:04Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 58 485 2011-09-01T15:22:50Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 59 486 2011-09-01T15:24:49Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 60 487 2011-09-01T15:27:00Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 61 491 2011-09-01T15:52:56Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 62 492 2011-09-01T15:53:26Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 63 493 2011-09-01T15:53:50Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 64 498 2011-09-01T15:59:56Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 65 499 2011-09-01T16:00:22Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 66 500 2011-09-01T16:00:46Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 67 501 2011-09-01T16:01:13Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 68 502 2011-09-01T16:02:25Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 69 503 2011-09-01T16:02:53Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 70 504 2011-09-01T16:03:18Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 56 505 483 2011-09-01T16:03:43Z Thunter 10 uploaded a new version of &quot;[[File:Rmsphase April2011.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 0 51 506 497 2011-09-01T16:06:37Z Thunter 10 wikitext text/x-wiki The following plots show the fraction of time suitable for observing in each EVLA band. The requirements for wind speed and API phase rms correspond to the values in the OPT. Click any plot to enlarge it. [[High_Frequency_Observing|Return to the High Frequency observing guide]] {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rx_January2011.png|center|250px]]'' '' | [[Image:rx_February2011.png|center|250px]]'' '' | [[Image:rx_March2011.png|center|250px]]'' '' |- ! April !! May !! June |- | [[Image:rx_April2011.png|center|250px]]'' '' | [[Image:rx_May2011.png|center|250px]]'' '' | [[Image:rx_June2011.png|center|250px]]'' '' |- ! July !! August !! September |- | [[Image:rx_July2011.png|center|250px]]'' '' | [[Image:rx_August2011.png|center|250px]]'' '' | [[Image:rx_September2010.png|center|250px]]'' '' |- ! October !! November !! December |- | [[Image:rx_October2010.png|center|250px]]'' '' | [[Image:rx_November2010.png|center|250px]]'' '' | [[Image:rx_December2010.png|center|250px]]'' '' |} 164e5d57f7cc626524e01ecd4d1a0f219578282a 521 506 2011-09-01T16:22:49Z Thunter 10 wikitext text/x-wiki The following plots show the fraction of time suitable for observing in each EVLA band. The simultaneous requirements for wind speed and API phase rms correspond to the values in the OPT. Click any plot to enlarge it. [[High_Frequency_Observing|Return to the High Frequency observing guide]] {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rx_January2011.png|center|250px]]'' '' | [[Image:rx_February2011.png|center|250px]]'' '' | [[Image:rx_March2011.png|center|250px]]'' '' |- ! April !! May !! June |- | [[Image:rx_April2011.png|center|250px]]'' '' | [[Image:rx_May2011.png|center|250px]]'' '' | [[Image:rx_June2011.png|center|250px]]'' '' |- ! July !! August !! September |- | [[Image:rx_July2011.png|center|250px]]'' '' | [[Image:rx_August2011.png|center|250px]]'' '' | [[Image:rx_September2010.png|center|250px]]'' '' |- ! October !! November !! December |- | [[Image:rx_October2010.png|center|250px]]'' '' | [[Image:rx_November2010.png|center|250px]]'' '' | [[Image:rx_December2010.png|center|250px]]'' '' |} 874f8a0a1d08f18a6761196eb7ae94e069e74193 523 521 2011-09-01T16:58:15Z Thunter 10 wikitext text/x-wiki The following plots show the fraction of time suitable for observing in each EVLA band for the one-year period September 2010 - August 2011. Click any plot to enlarge it. The simultaneous requirements for wind speed and API phase rms correspond to the current values in the OPT: {| border="1" class="wikitable" style="text-align: center;" |- ! Band: || 4, P, L || S || C || X || Ku || K || Ka || Q |- | API (deg) || any || <60 || <45 || <30 || <15 || <10 || <7 || <5 |- | wind (m/s) || any || any || any || <15 || <10 || <7 || <6 || <5 |} ! January !! February !! March [[High_Frequency_Observing|Return to the High Frequency observing guide]] {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rx_January2011.png|center|250px]]'' '' | [[Image:rx_February2011.png|center|250px]]'' '' | [[Image:rx_March2011.png|center|250px]]'' '' |- ! April !! May !! June |- | [[Image:rx_April2011.png|center|250px]]'' '' | [[Image:rx_May2011.png|center|250px]]'' '' | [[Image:rx_June2011.png|center|250px]]'' '' |- ! July !! August !! September |- | [[Image:rx_July2011.png|center|250px]]'' '' | [[Image:rx_August2011.png|center|250px]]'' '' | [[Image:rx_September2010.png|center|250px]]'' '' |- ! October !! November !! December |- | [[Image:rx_October2010.png|center|250px]]'' '' | [[Image:rx_November2010.png|center|250px]]'' '' | [[Image:rx_December2010.png|center|250px]]'' '' |} dcdeb9bc8b75df9cf1c2e83ff4e8e4c67b48551c 524 523 2011-09-01T16:58:28Z Thunter 10 wikitext text/x-wiki The following plots show the fraction of time suitable for observing in each EVLA band for the one-year period September 2010 - August 2011. Click any plot to enlarge it. The simultaneous requirements for wind speed and API phase rms correspond to the current values in the OPT: {| border="1" class="wikitable" style="text-align: center;" |- ! Band: || 4, P, L || S || C || X || Ku || K || Ka || Q |- | API (deg) || any || <60 || <45 || <30 || <15 || <10 || <7 || <5 |- | wind (m/s) || any || any || any || <15 || <10 || <7 || <6 || <5 |} [[High_Frequency_Observing|Return to the High Frequency observing guide]] {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rx_January2011.png|center|250px]]'' '' | [[Image:rx_February2011.png|center|250px]]'' '' | [[Image:rx_March2011.png|center|250px]]'' '' |- ! April !! May !! June |- | [[Image:rx_April2011.png|center|250px]]'' '' | [[Image:rx_May2011.png|center|250px]]'' '' | [[Image:rx_June2011.png|center|250px]]'' '' |- ! July !! August !! September |- | [[Image:rx_July2011.png|center|250px]]'' '' | [[Image:rx_August2011.png|center|250px]]'' '' | [[Image:rx_September2010.png|center|250px]]'' '' |- ! October !! November !! December |- | [[Image:rx_October2010.png|center|250px]]'' '' | [[Image:rx_November2010.png|center|250px]]'' '' | [[Image:rx_December2010.png|center|250px]]'' '' |} b11b621878e0f0d87178e50aa2342e48fa2a87d0 525 524 2011-09-01T17:00:16Z Thunter 10 wikitext text/x-wiki The following plots show the fraction of time suitable for observing in each EVLA band for the one-year period September 2010 - August 2011, according to the measured wind speed and API phase rms (opacity is not included). Click any plot to enlarge it. The simultaneous requirements for wind speed and API phase rms correspond to the current values in the OPT: {| border="1" class="wikitable" style="text-align: center;" |- ! Band: || 4, P, L || S || C || X || Ku || K || Ka || Q |- | API (deg) || any || <60 || <45 || <30 || <15 || <10 || <7 || <5 |- | wind (m/s) || any || any || any || <15 || <10 || <7 || <6 || <5 |} [[High_Frequency_Observing|Return to the High Frequency observing guide]] {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rx_January2011.png|center|250px]]'' '' | [[Image:rx_February2011.png|center|250px]]'' '' | [[Image:rx_March2011.png|center|250px]]'' '' |- ! April !! May !! June |- | [[Image:rx_April2011.png|center|250px]]'' '' | [[Image:rx_May2011.png|center|250px]]'' '' | [[Image:rx_June2011.png|center|250px]]'' '' |- ! July !! August !! September |- | [[Image:rx_July2011.png|center|250px]]'' '' | [[Image:rx_August2011.png|center|250px]]'' '' | [[Image:rx_September2010.png|center|250px]]'' '' |- ! October !! November !! December |- | [[Image:rx_October2010.png|center|250px]]'' '' | [[Image:rx_November2010.png|center|250px]]'' '' | [[Image:rx_December2010.png|center|250px]]'' '' |} 61e192e0b8b5e409014c221759ff3746a4853dd3 589 525 2011-09-02T00:03:06Z Thunter 10 wikitext text/x-wiki The following plots show the fraction of time that is most suitable for observing in each EVLA band for the one-year period September 2010 - August 2011, according to the measured wind speed and API phase rms (opacity is not included). Click any plot to enlarge it. The simultaneous requirements for wind speed and API phase rms correspond to the current values in the OPT: {| border="1" class="wikitable" style="text-align: center;" |- ! Band: || 4, P, L || S || C || X || Ku || K || Ka || Q |- | API (deg) || any || <60 || <45 || <30 || <15 || <10 || <7 || <5 |- | wind (m/s) || any || any || any || <15 || <10 || <7 || <6 || <5 |} [[High_Frequency_Observing|Return to the High Frequency observing guide]] {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rx_January2011.png|center|250px]]'' '' | [[Image:rx_February2011.png|center|250px]]'' '' | [[Image:rx_March2011.png|center|250px]]'' '' |- ! April !! May !! June |- | [[Image:rx_April2011.png|center|250px]]'' '' | [[Image:rx_May2011.png|center|250px]]'' '' | [[Image:rx_June2011.png|center|250px]]'' '' |- ! July !! August !! September |- | [[Image:rx_July2011.png|center|250px]]'' '' | [[Image:rx_August2011.png|center|250px]]'' '' | [[Image:rx_September2010.png|center|250px]]'' '' |- ! October !! November !! December |- | [[Image:rx_October2010.png|center|250px]]'' '' | [[Image:rx_November2010.png|center|250px]]'' '' | [[Image:rx_December2010.png|center|250px]]'' '' |} 0d6301c7f0cb32dadb6f9bc3879cfd7b25ed7f9d 591 589 2011-09-02T00:03:53Z Thunter 10 wikitext text/x-wiki The following plots show the fraction of time that was most suitable for observing in each EVLA band for the one-year period September 2010 - August 2011, according to the measured wind speed and API phase rms (opacity is not included). Click any plot to enlarge it. The simultaneous requirements for wind speed and API phase rms used to create these plots correspond to the current table of values in the OPT: {| border="1" class="wikitable" style="text-align: center;" |- ! Band: || 4, P, L || S || C || X || Ku || K || Ka || Q |- | API (deg) || any || <60 || <45 || <30 || <15 || <10 || <7 || <5 |- | wind (m/s) || any || any || any || <15 || <10 || <7 || <6 || <5 |} [[High_Frequency_Observing|Return to the High Frequency observing guide]] {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rx_January2011.png|center|250px]]'' '' | [[Image:rx_February2011.png|center|250px]]'' '' | [[Image:rx_March2011.png|center|250px]]'' '' |- ! April !! May !! June |- | [[Image:rx_April2011.png|center|250px]]'' '' | [[Image:rx_May2011.png|center|250px]]'' '' | [[Image:rx_June2011.png|center|250px]]'' '' |- ! July !! August !! September |- | [[Image:rx_July2011.png|center|250px]]'' '' | [[Image:rx_August2011.png|center|250px]]'' '' | [[Image:rx_September2010.png|center|250px]]'' '' |- ! October !! November !! December |- | [[Image:rx_October2010.png|center|250px]]'' '' | [[Image:rx_November2010.png|center|250px]]'' '' | [[Image:rx_December2010.png|center|250px]]'' '' |} 7077d76a19d9de21cb3e76a7530c4d5e13575192 593 591 2011-09-02T00:05:00Z Thunter 10 wikitext text/x-wiki The following plots show the fraction of time that was most suitable for observing in each EVLA band for the one-year period September 2010 - August 2011, according to the measured wind speed and API phase rms (opacity is not included). Click any plot to enlarge it. The simultaneous requirements for wind speed and API phase rms used to create these plots correspond to the current table of values in the OPT: {| border="1" class="wikitable" style="text-align: center;" |- ! Band: || 4, P, L || S || C || X || Ku || K || Ka || Q |- | API (deg) || any || <60 || <45 || <30 || <15 || <10 || <7 || <5 |- | wind (m/s) || any || any || any || <15 || <10 || <7 || <6 || <5 |} See also the [[Archival_API_Plots|monthly plots of API rms phase vs LST]] Return to the [[High_Frequency_Observing|High Frequency observing guide]] {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rx_January2011.png|center|250px]]'' '' | [[Image:rx_February2011.png|center|250px]]'' '' | [[Image:rx_March2011.png|center|250px]]'' '' |- ! April !! May !! June |- | [[Image:rx_April2011.png|center|250px]]'' '' | [[Image:rx_May2011.png|center|250px]]'' '' | [[Image:rx_June2011.png|center|250px]]'' '' |- ! July !! August !! September |- | [[Image:rx_July2011.png|center|250px]]'' '' | [[Image:rx_August2011.png|center|250px]]'' '' | [[Image:rx_September2010.png|center|250px]]'' '' |- ! October !! November !! December |- | [[Image:rx_October2010.png|center|250px]]'' '' | [[Image:rx_November2010.png|center|250px]]'' '' | [[Image:rx_December2010.png|center|250px]]'' '' |} 56da9e835f6bf96a2707f28d1b91085c7c5a6c41 594 593 2011-09-02T00:05:05Z Thunter 10 wikitext text/x-wiki The following plots show the fraction of time that was most suitable for observing in each EVLA band for the one-year period September 2010 - August 2011, according to the measured wind speed and API phase rms (opacity is not included). Click any plot to enlarge it. The simultaneous requirements for wind speed and API phase rms used to create these plots correspond to the current table of values in the OPT: {| border="1" class="wikitable" style="text-align: center;" |- ! Band: || 4, P, L || S || C || X || Ku || K || Ka || Q |- | API (deg) || any || <60 || <45 || <30 || <15 || <10 || <7 || <5 |- | wind (m/s) || any || any || any || <15 || <10 || <7 || <6 || <5 |} See also the [[Archival_API_Plots|monthly plots of API rms phase vs LST]] Return to the [[High_Frequency_Observing|High Frequency observing guide]] {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rx_January2011.png|center|250px]]'' '' | [[Image:rx_February2011.png|center|250px]]'' '' | [[Image:rx_March2011.png|center|250px]]'' '' |- ! April !! May !! June |- | [[Image:rx_April2011.png|center|250px]]'' '' | [[Image:rx_May2011.png|center|250px]]'' '' | [[Image:rx_June2011.png|center|250px]]'' '' |- ! July !! August !! September |- | [[Image:rx_July2011.png|center|250px]]'' '' | [[Image:rx_August2011.png|center|250px]]'' '' | [[Image:rx_September2010.png|center|250px]]'' '' |- ! October !! November !! December |- | [[Image:rx_October2010.png|center|250px]]'' '' | [[Image:rx_November2010.png|center|250px]]'' '' | [[Image:rx_December2010.png|center|250px]]'' '' |} 79b9c74c9d8a03841f088c6d2cda5e758b8e292f 6 57 508 484 2011-09-01T16:09:22Z Thunter 10 uploaded a new version of &quot;[[File:Rmsphase May2011.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 71 509 2011-09-01T16:10:11Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 72 510 2011-09-01T16:10:37Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 58 511 485 2011-09-01T16:11:14Z Thunter 10 uploaded a new version of &quot;[[File:Rmsphase June2011.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 0 49 512 490 2011-09-01T16:14:24Z Thunter 10 wikitext text/x-wiki These plots show the API rms phase vs LST for each month of the year. The magenta line in the top plot is the median computed per LST hour. Click any plot to enlarge it. Return to the [[Category:HighFrequency|High Frequency page]]. Return to the [[High_Frequency_Observing|High Frequency observing details page]]. {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rmsphase_January2011.png|center|250px]]'' '' | [[Image:rmsphase_February2011.png|center|250px]]'' '' | [[Image:rmsphase_March2011.png|center|250px]]'' '' |- ! April !! May !! June |- | [[Image:rmsphase_April2011.png|center|250px]]'' '' | [[Image:rmsphase_May2011.png|center|250px]]'' '' | [[Image:rmsphase_June2011.png|center|250px]]'' '' |- ! July !! August !! September |- | [[Image:rmsphase_July2011.png|center|250px]]'' '' | [[Image:rmsphase_August2011.png|center|250px]]'' '' | [[Image:rmsphase_September2010.png|center|250px]]'' '' |- ! October !! November !! December |- | [[Image:rmsphase_October2010.png|center|250px]]'' '' | [[Image:rmsphase_November2010.png|center|250px]]'' '' | [[Image:rmsphase_December2010.png|center|250px]]'' '' |} dc389ffd78d72fc0a7834e50ea98f020f2d49a06 513 512 2011-09-01T16:14:57Z Thunter 10 wikitext text/x-wiki These plots show the API rms phase vs LST for each month of the year. The magenta line in the top plot is the median computed per LST hour. Click any plot to enlarge it. Return to the [[HighFrequency|High Frequency page]]. Return to the [[High_Frequency_Observing|High Frequency observing details page]]. {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rmsphase_January2011.png|center|250px]]'' '' | [[Image:rmsphase_February2011.png|center|250px]]'' '' | [[Image:rmsphase_March2011.png|center|250px]]'' '' |- ! April !! May !! June |- | [[Image:rmsphase_April2011.png|center|250px]]'' '' | [[Image:rmsphase_May2011.png|center|250px]]'' '' | [[Image:rmsphase_June2011.png|center|250px]]'' '' |- ! July !! August !! September |- | [[Image:rmsphase_July2011.png|center|250px]]'' '' | [[Image:rmsphase_August2011.png|center|250px]]'' '' | [[Image:rmsphase_September2010.png|center|250px]]'' '' |- ! October !! November !! December |- | [[Image:rmsphase_October2010.png|center|250px]]'' '' | [[Image:rmsphase_November2010.png|center|250px]]'' '' | [[Image:rmsphase_December2010.png|center|250px]]'' '' |} b3958a5f26665938ccc1e94a80ebbe1cd165106a 515 513 2011-09-01T16:17:44Z Thunter 10 wikitext text/x-wiki These plots show the API rms phase vs LST for each month of the year. The magenta line in the top plot is the median computed per LST hour. Click any plot to enlarge it. Return to the [[High_Frequency_Observing|High Frequency observing details page]]. {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rmsphase_January2011.png|center|250px]]'' '' | [[Image:rmsphase_February2011.png|center|250px]]'' '' | [[Image:rmsphase_March2011.png|center|250px]]'' '' |- ! April !! May !! June |- | [[Image:rmsphase_April2011.png|center|250px]]'' '' | [[Image:rmsphase_May2011.png|center|250px]]'' '' | [[Image:rmsphase_June2011.png|center|250px]]'' '' |- ! July !! August !! September |- | [[Image:rmsphase_July2011.png|center|250px]]'' '' | [[Image:rmsphase_August2011.png|center|250px]]'' '' | [[Image:rmsphase_September2010.png|center|250px]]'' '' |- ! October !! November !! December |- | [[Image:rmsphase_October2010.png|center|250px]]'' '' | [[Image:rmsphase_November2010.png|center|250px]]'' '' | [[Image:rmsphase_December2010.png|center|250px]]'' '' |} 6be98a07816171e2288a6af8579bca5cd5d9d88f 595 515 2011-09-02T00:06:07Z Thunter 10 wikitext text/x-wiki These plots show the API rms phase vs LST for each month of the year. The magenta line in the top plot is the median computed per LST hour. Click any plot to enlarge it. See also the [[Monthly_Conditions_at_EVLA|monthly plots of the fraction of time available for each frequency band vs. LST]] Return to the [[High_Frequency_Observing|High Frequency observing details page]]. {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rmsphase_January2011.png|center|250px]]'' '' | [[Image:rmsphase_February2011.png|center|250px]]'' '' | [[Image:rmsphase_March2011.png|center|250px]]'' '' |- ! April !! May !! June |- | [[Image:rmsphase_April2011.png|center|250px]]'' '' | [[Image:rmsphase_May2011.png|center|250px]]'' '' | [[Image:rmsphase_June2011.png|center|250px]]'' '' |- ! July !! August !! September |- | [[Image:rmsphase_July2011.png|center|250px]]'' '' | [[Image:rmsphase_August2011.png|center|250px]]'' '' | [[Image:rmsphase_September2010.png|center|250px]]'' '' |- ! October !! November !! December |- | [[Image:rmsphase_October2010.png|center|250px]]'' '' | [[Image:rmsphase_November2010.png|center|250px]]'' '' | [[Image:rmsphase_December2010.png|center|250px]]'' '' |} 51d11143b38f5a77808095e87eff4ef1f99b89de 596 595 2011-09-02T00:06:57Z Thunter 10 wikitext text/x-wiki These plots show the API rms phase vs LST for each month of the year from September 2010 through August 2011. The magenta line in the top plot is the median computed per LST hour. Click any plot to enlarge it. See also the [[Monthly_Conditions_at_EVLA|monthly plots of the fraction of time available for each frequency band vs. LST]] Return to the [[High_Frequency_Observing|High Frequency observing details page]]. {| border="1" class="wikitable" style="text-align: center;" |- ! January !! February !! March |- | [[Image:rmsphase_January2011.png|center|250px]]'' '' | [[Image:rmsphase_February2011.png|center|250px]]'' '' | [[Image:rmsphase_March2011.png|center|250px]]'' '' |- ! April !! May !! June |- | [[Image:rmsphase_April2011.png|center|250px]]'' '' | [[Image:rmsphase_May2011.png|center|250px]]'' '' | [[Image:rmsphase_June2011.png|center|250px]]'' '' |- ! July !! August !! September |- | [[Image:rmsphase_July2011.png|center|250px]]'' '' | [[Image:rmsphase_August2011.png|center|250px]]'' '' | [[Image:rmsphase_September2010.png|center|250px]]'' '' |- ! October !! November !! December |- | [[Image:rmsphase_October2010.png|center|250px]]'' '' | [[Image:rmsphase_November2010.png|center|250px]]'' '' | [[Image:rmsphase_December2010.png|center|250px]]'' '' |} 459476c79e351640d4ec2d65b19600d7ff57c275 6 73 514 2011-09-01T16:17:35Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 59 516 486 2011-09-01T16:18:03Z Thunter 10 uploaded a new version of &quot;[[File:Rmsphase July2011.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 50 517 473 2011-09-01T16:19:07Z Thunter 10 uploaded a new version of &quot;[[File:Rmsphase August2011.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 52 518 476 2011-09-01T16:19:28Z Thunter 10 uploaded a new version of &quot;[[File:Rx August2011.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 60 519 487 2011-09-01T16:21:16Z Thunter 10 uploaded a new version of &quot;[[File:Rmsphase September2010.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 74 520 2011-09-01T16:21:35Z Thunter 10 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 64 522 498 2011-09-01T16:51:02Z Thunter 10 uploaded a new version of &quot;[[File:Rx January2011.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 75 563 2011-09-01T21:50:20Z Cbrogan 11 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 76 574 2011-09-01T23:45:30Z Cbrogan 11 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 10 2 610 255 2011-09-11T21:21:37Z Jmcmulli 2 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary '''Observational Status Summary'''] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:SpectraLine| Spectral Line Observations]] · [[:Category:Polarimetry| Polarimetry Observations]] · [[:Category:Planetary| Planetary Observations]] · [[:Category:HighFrequency| High Frequency Observing (K, Ka, Q)]] · [[:Category:LowFrequency| Low Frequency Observing (L, S, C)]] · [[:Category:PhasedArray| Phased Array Observing]] [[:Category:Pulsar| Pulsar Observing]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] · [[Key to Calcodes]] |} 5399b1e21c9a93a19be79ee7be28eb2638d76324 872 610 2012-01-25T18:26:03Z Dshepher 15 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary '''Observational Status Summary'''] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:SpectraLine| Spectral Line Observations]] · [[:Category:Polarimetry| Polarimetry Observations]] · [[:Category:Planetary| Planetary Observations]] · [[:Category:HighFrequency| High Frequency Observing (K, Ka, Q)]] · [[:Category:LowFrequency| Low Frequency Observing (L, S, C)]] · [[:Category:PhasedArray| Phased Array Observing]] [[:Category:Pulsar| Pulsar Observing]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] · [[Key to Calcodes]] {|style="background-color:#EAF5FB;" | [[:Category:Post-Processing | EVLA Reduction Strategy | Special Considerations for EVLA Data Calibration and Imaging in AIPS ]] | [[Image:OrionA_Kspectrum.png|thumb]] |} |} 9820cd99f0d780764a1af4527cfa4d2a170a3ab4 873 872 2012-01-25T18:28:18Z Dshepher 15 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary '''Observational Status Summary'''] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:SpectraLine| Spectral Line Observations]] · [[:Category:Polarimetry| Polarimetry Observations]] · [[:Category:Planetary| Planetary Observations]] · [[:Category:HighFrequency| High Frequency Observing (K, Ka, Q)]] · [[:Category:LowFrequency| Low Frequency Observing (L, S, C)]] · [[:Category:PhasedArray| Phased Array Observing]] [[:Category:Pulsar| Pulsar Observing]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] · [[Key to Calcodes in AIPS]] · [[:Category:Post-Processing | EVLA Reduction Strategy | Special Considerations for EVLA Data Calibration and Imaging in AIPS]] |} e79017acf5d6f27e8c1aa2837cc012b6438c34e7 0 48 611 422 2011-09-22T17:54:29Z Jott 8 wikitext text/x-wiki spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction 1378636b26a2ead1d33efbb763f450bfe56502f7 612 611 2011-09-22T17:56:10Z Jott 8 wikitext text/x-wiki spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 28b05d8270e74ced23b470bfead4d563654b2f3d 614 612 2011-10-04T11:24:04Z Jott 8 wikitext text/x-wiki spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 3f439fe5e27d41494aa8f3e0642db16c43650efc 14 21 613 279 2011-10-03T20:51:19Z Smyers 4 /* Observing Recommendations */ wikitext text/x-wiki = Polarization Calibration = == Revised OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Detailed Guidelines == === Observing Recommendations === There are several strategies for deriving the Q/U angle calibration: * Observation of a primary polarization standard (Category A) * Observation of a secondary polarization calibrator (Category B with Note 3) with auxilary monitoring observations to transfer from primary. This calibration is needed to set the polarization vector angle 0.5*arctan(U/Q) and should be done in all cases. There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (Category C) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source (Category B or your gain calibrator, if bright enough) * Two scans of a source of known polarization (Category A or B with transfer) See Tables 1-4 below for Category A-D source catalogs. ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Wide Band Considerations ==== * Performance over 1, 2 and 8 GHz band widths ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source (or your gain calibrator) in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (>4GHz) || A1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || A2 |- |} Notes: *A1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *A2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || B1 |- | J0555+3948 || B0552+398 || bright, flat spectrum, monitored, moderate variability || B1,B2 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || B1 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || B1,B2 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || B1 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || B3 |- |} Notes: *B1. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *B2. Low polarization at low frequencies (L, sometimes S,C), do not use as angle calibrator. *B3. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || C1 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || C2 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || C3 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || C2 |- |} Notes: *C1. Very bright and low polarization (<1%), but variable flux density. Approaches 1% polarized at 43GHz. *C2. Weak at high frequency, but stable flux and very low polarization. *C3. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === For CASA reduction and analysis of polarization data, please see the following links: * [[http://casa.nrao.edu/docs/UserMan/UserMansu184.html#x213-2100004.4.5| Instrumental Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu185.html#x214-2110004.4.5.1| Heuristics and Strategies for Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu186.html#x215-2120004.4.5.2| A Note on Channelized Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu187.html#x216-2130004.4.5.3| A Polarization Calibration Example]] 690848f88ff2244fc9dbd0b1dea2edc7418afab4 615 613 2011-10-04T14:57:56Z Smyers 4 /* Revised OSS Guidelines */ wikitext text/x-wiki = Polarization Calibration = == Revised OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 7.5 cm or shorter (> 4GHz). The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Detailed Guidelines == === Observing Recommendations === There are several strategies for deriving the Q/U angle calibration: * Observation of a primary polarization standard (Category A) * Observation of a secondary polarization calibrator (Category B with Note 3) with auxilary monitoring observations to transfer from primary. This calibration is needed to set the polarization vector angle 0.5*arctan(U/Q) and should be done in all cases. There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (Category C) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source (Category B or your gain calibrator, if bright enough) * Two scans of a source of known polarization (Category A or B with transfer) See Tables 1-4 below for Category A-D source catalogs. ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Wide Band Considerations ==== * Performance over 1, 2 and 8 GHz band widths ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source (or your gain calibrator) in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (>4GHz) || A1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || A2 |- |} Notes: *A1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *A2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || B1 |- | J0555+3948 || B0552+398 || bright, flat spectrum, monitored, moderate variability || B1,B2 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || B1 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || B1,B2 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || B1 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || B3 |- |} Notes: *B1. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *B2. Low polarization at low frequencies (L, sometimes S,C), do not use as angle calibrator. *B3. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || C1 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || C2 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || C3 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || C2 |- |} Notes: *C1. Very bright and low polarization (<1%), but variable flux density. Approaches 1% polarized at 43GHz. *C2. Weak at high frequency, but stable flux and very low polarization. *C3. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === For CASA reduction and analysis of polarization data, please see the following links: * [[http://casa.nrao.edu/docs/UserMan/UserMansu184.html#x213-2100004.4.5| Instrumental Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu185.html#x214-2110004.4.5.1| Heuristics and Strategies for Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu186.html#x215-2120004.4.5.2| A Note on Channelized Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu187.html#x216-2130004.4.5.3| A Polarization Calibration Example]] 2f8f7a8e5a1868c90e92b87a07d6419637f22b90 616 615 2011-10-04T14:58:20Z Smyers 4 /* Revised OSS Guidelines */ wikitext text/x-wiki = Polarization Calibration = == Revised OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 7.5 cm or shorter (4GHz and higher). The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Detailed Guidelines == === Observing Recommendations === There are several strategies for deriving the Q/U angle calibration: * Observation of a primary polarization standard (Category A) * Observation of a secondary polarization calibrator (Category B with Note 3) with auxilary monitoring observations to transfer from primary. This calibration is needed to set the polarization vector angle 0.5*arctan(U/Q) and should be done in all cases. There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (Category C) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source (Category B or your gain calibrator, if bright enough) * Two scans of a source of known polarization (Category A or B with transfer) See Tables 1-4 below for Category A-D source catalogs. ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Wide Band Considerations ==== * Performance over 1, 2 and 8 GHz band widths ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source (or your gain calibrator) in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (>4GHz) || A1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || A2 |- |} Notes: *A1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *A2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || B1 |- | J0555+3948 || B0552+398 || bright, flat spectrum, monitored, moderate variability || B1,B2 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || B1 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || B1,B2 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || B1 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || B3 |- |} Notes: *B1. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *B2. Low polarization at low frequencies (L, sometimes S,C), do not use as angle calibrator. *B3. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || C1 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || C2 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || C3 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || C2 |- |} Notes: *C1. Very bright and low polarization (<1%), but variable flux density. Approaches 1% polarized at 43GHz. *C2. Weak at high frequency, but stable flux and very low polarization. *C3. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === For CASA reduction and analysis of polarization data, please see the following links: * [[http://casa.nrao.edu/docs/UserMan/UserMansu184.html#x213-2100004.4.5| Instrumental Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu185.html#x214-2110004.4.5.1| Heuristics and Strategies for Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu186.html#x215-2120004.4.5.2| A Note on Channelized Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu187.html#x216-2130004.4.5.3| A Polarization Calibration Example]] afbbf17fca9d9998c117d6e43be5ab65ff60695f 617 616 2011-10-04T15:00:34Z Smyers 4 /* Revised OSS Guidelines */ wikitext text/x-wiki = Polarization Calibration = == Revised OSS Guidelines == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). This can usually be obtained for your gain calibrator if it is at Dec>20 and your track is 4 hours or longer - check the reporting from the OPT when setting up your observations. If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 7.5 cm or shorter (4GHz and higher). The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Detailed Guidelines == === Observing Recommendations === There are several strategies for deriving the Q/U angle calibration: * Observation of a primary polarization standard (Category A) * Observation of a secondary polarization calibrator (Category B with Note 3) with auxilary monitoring observations to transfer from primary. This calibration is needed to set the polarization vector angle 0.5*arctan(U/Q) and should be done in all cases. There are several strategies for deriving the instrumental polarization: * Single scan observation of a zero polarization source (Category C) * Several scans (minimum of 3 over 60 degrees of parallactic angle) of an unknown polarization source (Category B or your gain calibrator, if bright enough) * Two scans of a source of known polarization (Category A or B with transfer) See Tables 1-4 below for Category A-D source catalogs. ==== Low Frequency Considerations ==== * Ionosphere monitoring Global Ionospheric TEC maps are available via: http://iono.jpl.nasa.gov/latest_rti_global.html * Solar activity monitoring Solar activity and general space weather can be reviewed at the NOAA site: http://www.swpc.noaa.gov/today.html The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. ==== High Frequency Considerations ==== * Source list restrictions * Need longer observations ==== Wide Band Considerations ==== * Performance over 1, 2 and 8 GHz band widths ==== Time stability ==== Current results (by band) ==== Frequency stability ==== Current results (by band) ==== Polarization Calibrator Catalog and Selection ==== The following sources are known to be useable for polarization calibration. These consist of a few "pol standard" sources with known stable polarization (for Q/U angle calibration), plus a number of "bright" sources with "monitored" variable flux densities and polarization. Some of these are seen to have only "moderate variablity" and could be used as secondary angle calibrators if you can transfer the angle from the monitoring observations. Assume others (particularly "flat spectrum") are highly variable. There are also a few "bright, low pol" sources available as leakage calibrators (but the can have measureable polarization at high frequencies). NOTE: Be sure to use the EVLA OPT Source catalog to get the standard J2000 positions and see the fluxes. Calibration Selection Procedure: * Select Polarization Standard (to calibrate polarization angle Q/U) - optimally select one Category A source and observe at least one scan. Alternative: use a "moderately variable" Category B calibrator and use monitoring information (may have to submit your own SB for this) to transfer from a primary. * Select Leakage Calibrator (to determine intrumental polarization) - optimally select one Category C low-polarization source or Category B secondary source (or your gain calibrator) in optimal Dec range (see Table 2 note 3) for PA coverage during run (if long enough). Single scans ok for Category C. Alternative: try a Category D CSO if no other options available. {| border="1" align="center" |+ '''Table 1: Category A - primary polarization standards''' !Source !Other name !Comments !Notes |- | J0137+3309 || B0134+329 (3C48) || pol standard (>4GHz) || A1 |- | J0521+1638 || B0518+165 (3C138) || pol standard || |- | J1331+3030 || B1328+307 (3C286) || pol standard || A2 |- |} Notes: *A1. 3C48 is weak at high frequency and somewhat resolved in larger configurations. Depolarized below 4GHz. *A2. 3C286 is our foremost primary calibrator and should be used if available. {| border="1" align="center" |+ '''Table 2: Category B - secondary polarization calibrators''' !Source !Other name !Comments !Notes |- | J0359+5057 || B0355+508 (NRAO150) || bright, flat spectrum, monitored, moderate variability || B1 |- | J0555+3948 || B0552+398 || bright, flat spectrum, monitored, moderate variability || B1,B2 |- | J0854+2006 || B0851+202 || bright, flat spectrum, monitored, moderate variability || B1 |- | J0927+3902 || B0923+392 || bright, flat spectrum, monitored, moderate variability || B1,B2 |- | J1310+3220 || B1308+326 || monitored || |- | J2136+0041 || B2134+004 || bright, flat spectrum, monitored, moderate variability || |- | J2202+4216 || B2200+420 (BLLac) || bright, flat spectrum, monitored, moderate variability || B1 |- | J2253+1608 || B2251+158 (3C454.3) || bright, flat spectrum, monitored || B3 |- |} Notes: *B1. In optimal Declination range to be used as leakage calibrator with PA coverage. Recommended as calibrators and if necessary can be used as secondary standards with monitoring. *B2. Low polarization at low frequencies (L, sometimes S,C), do not use as angle calibrator. *B3. Highly variable and interesting in its own right. {| border="1" align="center" |+ '''Table 3: Category C - primary low polarization leakage calibrators''' !Source !Other name !Comments !Notes |- | J0319+4130 || B0316+413 (3C84) || low pol, bright, flat spectrum, monitored || C1 |- | J0713+4349 || B0710+439 || low pol, CSO, monitored || C2 |- | J1407+2827 || B1404+286 (OQ208) || low pol, steep spectrum || C3 |- | J2355+4950 || B2352+495 || low pol, CSO, monitored || C2 |- |} Notes: *C1. Very bright and low polarization (<1%), but variable flux density. Approaches 1% polarized at 43GHz. *C2. Weak at high frequency, but stable flux and very low polarization. *C3. Weak at high frequency, bright and low polarization below 9GHz. The following northern sources are known to be CSO (Compact Symmetric Objects) and are characteristically unpolarized. They can be used over a range of frequencies (Gugliucci, N.E. et al. 2007, ApJ 661, 78) as "low pol" leakage calibrators. CSOs tend to be on the weak side and should be used with care at higher frequencies. We have not used these with the EVLA and thus rate them as "secondary" unpolarized calibrators. Let us know if you use these so we can evaluate their performance. WARNING: the positions given below are B1950, use the Source names in the EVLA OPT to get the J2000 positions. {| border="1" align="center" |+ '''Table 4: Category D - secondary (unverified) low polarization sources''' !Source !RA (1950) !DEC (1950) !B1950 Name' !Comments |- | J0029+3456 || 00 26 34.8386 || 34 39 57.586 || 0026+346 || CSO |- | J0111+3906 || 01 08 47.2595 || 38 50 32.691 || 0108+388 || CSO |- | J0410+7656 || 04 03 58.60 || 76 48 54.0 || 0404+768 || CSO |- | J1035+5628 || 10 31 55.9562 || 56 44 18.284 || 1031+567 || CSO |- | J1148+5924 || 11 46 10.4160 || 59 41 36.834 || 1146+596 || CSO |- | J1400+6210 || 13 58 58.310 || 62 25 08.40 || 1358+624 || CSO |- | J1815+6127 || 18 15 05.4851 || 61 26 04.496 || 1815+614 || CSO |- | J1823+7938 || 18 26 43.2676 || 79 36 59.943 || 1826+796 || CSO |- | J1944+5448 || 19 43 22.6729 || 54 40 47.955 || 1943+546 || CSO |- | J1945+7055 || 19 46 12.0492 || 70 48 21.397 || 1946+708 || CSO |- | J2022+6136 || 20 21 13.3005 || 61 27 18.157 || 2021+614 || CSO |- |} Comments: * at least one "pol standard" (ideally from Category A) should be included for angle calibration * "bright" sources are easily useable as leakage calibrators with PA coverage (and probably good for bandpasses to boot!) * "monitored" sources can be found at http://www.vla.nrao.edu/astro/calib/polar/ * "steep spectrum" sources are likely weak at high frequencies * "flat spectrum" sources are likely bright at high frequencies but variable * "moderately variable" sources may be useable in a pinch if you can get a nearby (in time) monitoring observation (see below) ==== Monitoring Observations ==== For the VLA, a decade-long monitoring program was carried out with the goal of allowing transfer from our standard sources to bright sources useable as VLBA calibrators. The results of this can be found at http://www.vla.nrao.edu/astro/calib/polar/ We are in the process of beginning such a program for the EVLA. There is no pipeline produced monitoring results as of this time, but intrepid users can find the data in the public archive https://archive.nrao.edu/archive/archiveproject.jsp under project code TPOL0003. The VLA database (particularly before the transition in 2008) can be used to see the level of variability in these sources, and to get an idea of the flux density ranges to expect. === Post-processing Guidelines === For CASA reduction and analysis of polarization data, please see the following links: * [[http://casa.nrao.edu/docs/UserMan/UserMansu184.html#x213-2100004.4.5| Instrumental Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu185.html#x214-2110004.4.5.1| Heuristics and Strategies for Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu186.html#x215-2120004.4.5.2| A Note on Channelized Polarization Calibration]] ** [[http://casa.nrao.edu/docs/UserMan/UserMansu187.html#x216-2130004.4.5.3| A Polarization Calibration Example]] 70fb3186b0a91266bb4d6c5f1af86462877ccf55 6 78 642 2011-11-17T22:07:02Z Twiegert 13 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 79 646 2011-11-17T22:12:58Z Twiegert 13 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 80 719 2011-12-15T21:28:59Z Gvanmoor 7 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 0 1 862 225 2012-01-24T23:43:33Z Dshepher 15 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="99%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">Featured Article</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{Featured Article}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="1%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">News</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | * 02-Mar-2010: First science observations with WIDAR * 12-Apr-2010: First 27-antenna correlation * 25-Jun-2010: First fringes with 3-bit samplers (3 antennas) * 05-Jul-2010: First 2 GHz BW science observing * 09-Sep-2010: First 74 MHz observations with WIDAR * 13-Oct-2010: Time averaging enabled in correlator back-end {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:left;" | <div style="font-size:120%">Configuration & Proposal Dates</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| |- | style="padding-left:6px; padding-top:6px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| ! Trimester ! Observing Period ! Configuration ! Proposal Deadline |- | 2010 Mar 01 - 2010 Sep 13 | align="center" | D | 2009 Oct 1 |- | 2010 Sep 17 - 2010 Oct 04 | align="center" | DnC | 2009 Oct 1 |- | 2010 Oct 15 - 2011 Jan 18 | align="center" | C | 2010 Feb 1 |- | 2011 Jan 21 - 2011 Feb 07 | align="center" | CnB | 2010 Feb 1 |- | 2011 Feb 18 - 2011 May 09 | align="center" | B | 2010 Jun 1 |- | 2011 May 13 - 2011 May 31 | align="center" | BnA | 2010 Jun 1 |- | 2011 Jun 10 - 2011 Sep 12 | align='center' | A | 2010 Oct 1 |} <!-- TEXT 2 (END) --> |} |} f8a7bb62b4f554d39994dd08d895fe6beb67c5d3 863 862 2012-01-24T23:44:19Z Dshepher 15 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">Featured Article</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{Featured Article}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">News</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | * 02-Mar-2010: First science observations with WIDAR * 12-Apr-2010: First 27-antenna correlation * 25-Jun-2010: First fringes with 3-bit samplers (3 antennas) * 05-Jul-2010: First 2 GHz BW science observing * 09-Sep-2010: First 74 MHz observations with WIDAR * 13-Oct-2010: Time averaging enabled in correlator back-end {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:left;" | <div style="font-size:120%">Configuration & Proposal Dates</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| |- | style="padding-left:6px; padding-top:6px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| ! Trimester ! Observing Period ! Configuration ! Proposal Deadline |- | 2010 Mar 01 - 2010 Sep 13 | align="center" | D | 2009 Oct 1 |- | 2010 Sep 17 - 2010 Oct 04 | align="center" | DnC | 2009 Oct 1 |- | 2010 Oct 15 - 2011 Jan 18 | align="center" | C | 2010 Feb 1 |- | 2011 Jan 21 - 2011 Feb 07 | align="center" | CnB | 2010 Feb 1 |- | 2011 Feb 18 - 2011 May 09 | align="center" | B | 2010 Jun 1 |- | 2011 May 13 - 2011 May 31 | align="center" | BnA | 2010 Jun 1 |- | 2011 Jun 10 - 2011 Sep 12 | align='center' | A | 2010 Oct 1 |} <!-- TEXT 2 (END) --> |} |} 1dda164604a57accd6ca4d450320e45232e2282a 864 863 2012-01-24T23:51:10Z Dshepher 15 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">Featured Article</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{Featured Article}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">News</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | * 02-Mar-2010: First science observations with WIDAR * 12-Apr-2010: First 27-antenna correlation * 25-Jun-2010: First fringes with 3-bit samplers (3 antennas) * 05-Jul-2010: First 2 GHz BW science observing * 09-Sep-2010: First 74 MHz observations with WIDAR * 13-Oct-2010: Time averaging enabled in correlator back-end {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #ddddc0; text-align:left;" | <div style="font-size:120%">Configuration & Proposal Dates</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| |- | style="padding-left:6px; padding-top:6px; padding-bottom:2px; padding-right:0px; font-size:8pt" | {| |} <!-- TEXT 2 (END) --> |} |} 996896396229cc1eb43e3b0c5f5079139ec9bea4 865 864 2012-01-24T23:52:21Z Dshepher 15 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 2 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">Featured Article</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{Featured Article}} <!-- TEXT 2 (END) --> |} <!-- BLOCK 2 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">News</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | * 02-Mar-2010: First science observations with WIDAR * 12-Apr-2010: First 27-antenna correlation * 25-Jun-2010: First fringes with 3-bit samplers (3 antennas) * 05-Jul-2010: First 2 GHz BW science observing * 09-Sep-2010: First 74 MHz observations with WIDAR * 13-Oct-2010: Time averaging enabled in correlator back-end {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- |} bc77abe2e0623b3cd5a42ee58ef3445e2de2aaca 14 27 866 236 2012-01-25T00:00:23Z Dshepher 15 wikitext text/x-wiki = Low Frequency Observing (L, S, C) = ==Introduction== This document is intended for observers planning EVLA observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the EVLA, consult the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for EVLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst EVLA Proposal Submission Tool (PST)]. ===Correlator Setup=== Starting with the D-configuration in 2011 OSRO modes can use up to 8 contiguous subbands per intermediate frequency passband ("baseband") with a maximum (minimum) bandwidth of 128 MHz (0.03 MHz) per subband. All subbands must have the same width and number of channels. There are 2 such independently tunable basebands available. See the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page] for a complete description. You should pick the correlator mode that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Low_Frequency_Observing#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see for example [http://en.wikipedia.org/wiki/Gibbs_phenomenon]). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Observing Frequencies=== All EVLA antennas have been equipped with L and C band receivers. As of January 2012, 20 EVLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Sensitivity Observational status summary page]. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Low_Frequency_Observing#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the EVLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider #what angular resolution is required for your science at the desired observing frequency #for resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission #how much flux will be resolved out by the array configuration that gives the desired angular resolution #will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Low_Frequency_Observing#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * To assess if the calibrator is appropriate for the selected bandwidth, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq 4 </math> within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * To assess if the bandpass calibrator is appropriate for the selected channel width, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the the S/N is <math> \geq 4 </math> within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] ac1aded25c23635611a0b762a95ac3c9d55bc380 867 866 2012-01-25T17:52:24Z Dshepher 15 /* Introduction */ wikitext text/x-wiki = Low Frequency Observing (L, S, C) = ==Introduction== This document is intended for observers planning EVLA observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the EVLA, consult the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for EVLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst EVLA Proposal Submission Tool (PST)]. ===Correlator Setup=== Starting with the D-configuration in 2011 OSRO modes can use up to 8 contiguous subbands per intermediate frequency passband ("baseband") with a maximum (minimum) bandwidth of 128 MHz (0.03 MHz) per subband. All subbands must have the same width and number of channels. There are 2 such independently tunable basebands available. See the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page] for a complete description. You should pick the correlator mode that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Low_Frequency_Observing#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see for example [http://en.wikipedia.org/wiki/Gibbs_phenomenon]). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Observing Frequencies=== All EVLA antennas have been equipped with L and C band receivers. As of January 2012, 20 EVLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Sensitivity Observational status summary page]. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Low_Frequency_Observing#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the EVLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider #what angular resolution is required for your science at the desired observing frequency #for resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission #how much flux will be resolved out by the array configuration that gives the desired angular resolution #will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Low_Frequency_Observing#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * To assess if the calibrator is appropriate for the selected bandwidth, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq 4 </math> within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * To assess if the bandpass calibrator is appropriate for the selected channel width, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the the S/N is <math> \geq 4 </math> within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 7bf9f1dc7b50292487086f69241aa339387ee529 869 867 2012-01-25T18:10:51Z Dshepher 15 wikitext text/x-wiki = Low Frequency Observing (L, S, C) = ==Introduction== This document is intended for observers planning EVLA observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the EVLA, consult the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for EVLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst EVLA Proposal Submission Tool (PST)]. ===Correlator Setup=== Starting with the D-configuration in 2011 OSRO modes can use up to 8 contiguous subbands per intermediate frequency passband ("baseband") with a maximum (minimum) bandwidth of 128 MHz (0.03 MHz) per subband. All subbands must have the same width and number of channels. There are 2 such independently tunable basebands available. See the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page] for a complete description. You should pick the correlator mode that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see for example [http://en.wikipedia.org/wiki/Gibbs_phenomenon]). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Observing Frequencies=== All EVLA antennas have been equipped with L and C band receivers. As of January 2012, 20 EVLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Sensitivity Observational status summary page]. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Low_Frequency_Observing#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the EVLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider #what angular resolution is required for your science at the desired observing frequency #for resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission #how much flux will be resolved out by the array configuration that gives the desired angular resolution #will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Low_Frequency_Observing#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * To assess if the calibrator is appropriate for the selected bandwidth, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq 4 </math> within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * To assess if the bandpass calibrator is appropriate for the selected channel width, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the the S/N is <math> \geq 4 </math> within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 9ba6b64f1280bc8c3a9669ef2f6af3d8ed1ce3bf 871 869 2012-01-25T18:12:38Z Dshepher 15 wikitext text/x-wiki = Low Frequency Observing (L, S, C) = ==Introduction== This document is intended for observers planning EVLA observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the EVLA, consult the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for EVLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst EVLA Proposal Submission Tool (PST)]. ===Correlator Setup=== Starting with the D-configuration in 2011 OSRO modes can use up to 8 contiguous subbands per intermediate frequency passband ("baseband") with a maximum (minimum) bandwidth of 128 MHz (0.03 MHz) per subband. All subbands must have the same width and number of channels. There are 2 such independently tunable basebands available. See the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page] for a complete description. You should pick the correlator mode that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see for example [http://en.wikipedia.org/wiki/Gibbs_phenomenon]). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Observing Frequencies=== All EVLA antennas have been equipped with L and C band receivers. As of January 2012, 20 EVLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Sensitivity Observational status summary page]. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the EVLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider #what angular resolution is required for your science at the desired observing frequency #for resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission #how much flux will be resolved out by the array configuration that gives the desired angular resolution #will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * To assess if the calibrator is appropriate for the selected bandwidth, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq 4 </math> within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * To assess if the bandpass calibrator is appropriate for the selected channel width, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the the S/N is <math> \geq 4 </math> within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] f2495c1d99b81fe3d239b4af4e52e45bd172871f 10 2 874 873 2012-01-25T18:28:42Z Dshepher 15 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary '''Observational Status Summary'''] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:SpectraLine| Spectral Line Observations]] · [[:Category:Polarimetry| Polarimetry Observations]] · [[:Category:Planetary| Planetary Observations]] · [[:Category:HighFrequency| High Frequency Observing (K, Ka, Q)]] · [[:Category:LowFrequency| Low Frequency Observing (L, S, C)]] · [[:Category:PhasedArray| Phased Array Observing]] [[:Category:Pulsar| Pulsar Observing]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> [[:Category:Post-Processing|EVLA Reduction Strategy]] · [http://casaguides.nrao.edu '''CASA Reduction Guides'''] · [[Key to Calcodes]] · [[:Category:Post-Processing | EVLA Reduction Strategy | Special Considerations for EVLA Data Calibration and Imaging in AIPS]] |} 64a35db18118481f942e5b2512603e00a274fa83 875 874 2012-01-25T18:32:37Z Dshepher 15 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary '''Observational Status Summary'''] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:SpectraLine| Spectral Line Observations]] · [[:Category:Polarimetry| Polarimetry Observations]] · [[:Category:Planetary| Planetary Observations]] · [[:Category:HighFrequency| High Frequency Observing (K, Ka, Q)]] · [[:Category:LowFrequency| Low Frequency Observing (L, S, C)]] · [[:Category:PhasedArray| Phased Array Observing]] [[:Category:Pulsar| Pulsar Observing]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> CASA: [http://casaguides.nrao.edu '''CASA Reduction Guides'''] <br> AIPS: [[Key to Calcodes]] · [[:Category:Post-Processing|EVLA Reduction Strategy | Special Considerations for EVLA Data Calibration and Imaging in AIPS]]<br> |} 6bf1753c17ab8a37e3bd208f9f50a8da16521b5d 876 875 2012-01-25T18:34:00Z Dshepher 15 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary '''Observational Status Summary'''] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:SpectraLine| Spectral Line Observations]] · [[:Category:Polarimetry| Polarimetry Observations]] · [[:Category:Planetary| Planetary Observations]] · [[:Category:HighFrequency| High Frequency Observing (K, Ka, Q)]] · [[:Category:LowFrequency| Low Frequency Observing (L, S, C)]] · [[:Category:PhasedArray| Phased Array Observing]] [[:Category:Pulsar| Pulsar Observing]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> CASA: [http://casaguides.nrao.edu '''CASA Reduction Guides'''] <br> AIPS: [[Key to Calcodes]] · [[:Category:Post-Processing|Special Considerations for EVLA Data Calibration and Imaging in AIPS]]<br> |} e3274eabe8e699142f0f196e90e80eab8f73e726 0 1 877 865 2012-01-25T18:37:51Z Dshepher 15 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">News</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | * 02-Mar-2010: First science observations with WIDAR * 12-Apr-2010: First 27-antenna correlation * 25-Jun-2010: First fringes with 3-bit samplers (3 antennas) * 05-Jul-2010: First 2 GHz BW science observing * 09-Sep-2010: First 74 MHz observations with WIDAR * 13-Oct-2010: Time averaging enabled in correlator back-end {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- |} b7a887cb605959301da0ed517ba5bebb25524915 878 877 2012-01-25T18:39:44Z Dshepher 15 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT BOX ************************** --> |width="52%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 1 (END) --> <!-- *********************** RIGHT BOX **************************** --> |width="48%" style="background:#fffff3; border:1px solid #eeeed1; font-size:100%; -moz-border-radius-topright:0px; -moz-border-radius-bottomright:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 2: Featured Article --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- ! style="background:#eeeed1; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">News</div> <!-- TITLE 2 (END)--> <!-- TEXT 2 --> |- | style="background:#fffff3; padding-left:0px; padding-top:2px; padding-bottom:2px; padding-right:0px; font-size:8pt" | * 02-Mar-2010: First science observations with WIDAR * 12-Apr-2010: First 27-antenna correlation * 25-Jun-2010: First fringes with 3-bit samplers (3 antennas) * 05-Jul-2010: First 2 GHz BW science observing * 09-Sep-2010: First 74 MHz observations with WIDAR * 13-Oct-2010: Time averaging enabled in correlator back-end {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#fffff3;" <!-- TITLE 2 --> |- |} b77edc9bb50cab9645cda5496f478c6d8392c853 879 878 2012-01-25T18:41:48Z Dshepher 15 wikitext text/x-wiki <!-- BANNER ACROSS TOP OF PAGE --> {| id="mp-topbanner" style="width:100%; background:#fcfcfc; margin-top:1.2em; border:1px solid #ccc;" | style="width:56%; color:#000;" | <!-- "WELCOME TO EVLAGUIDES" --> {| style="width:280px; border:none; background:none;" | style="width:280px; text-align:center; white-space:nowrap; color:#000;" | <div style="font-size:162%; border:none; margin:0; padding:.1em; color:#000;">Welcome to EVLA Guides [[File:vla_panorama_lo.jpg|300px|center]] </div> |} {|width="100%" cellspacing="10" cellpadding="0" |- <!-- ******************** LEFT & ONLY BOX ************************** --> |width="100%" style="background:#EAF5FB; border:1px solid #e1eaee; font-size:100%; -moz-border-radius-topleft:0px; -moz-border-radius-bottomleft:0px; padding:7px 7px 7px 7px;" valign="top"| <!-- BLOCK 1: EVLA Guides --> {| width=100% cellpadding="0" cellspacing="0" valign="top" style="background:#F1FAFF;" <!-- TITLE 1 --> |- ! style="background:#e1eaee; border:1px solid #d0d9dd; text-align:left" | <div style="font-size:120%">EVLA Guides</div> <!-- TITLE 1 (END)--> <!-- TEXT 1 --> |- | style="background:#EAF5FB; padding-left:0px; padding-top:2px; padding-bottom:2px;" | {{EVLA Guides}} <!-- TEXT 1 (END) --> |} <!-- BLOCK 1 (END) --> 7257550f08ff5edda86a0d9cd08628e807c1c350 14 27 880 871 2012-01-25T21:19:42Z Dshepher 15 wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S, C)== This document is intended for observers planning EVLA observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the EVLA, consult the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for EVLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst EVLA Proposal Submission Tool (PST)]. ===Correlator Setup=== Starting with the D-configuration in 2011 OSRO modes can use up to 8 contiguous subbands per intermediate frequency passband ("baseband") with a maximum (minimum) bandwidth of 128 MHz (0.03 MHz) per subband. All subbands must have the same width and number of channels. There are 2 such independently tunable basebands available. See the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page] for a complete description. You should pick the correlator mode that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see for example [http://en.wikipedia.org/wiki/Gibbs_phenomenon]). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Observing Frequencies=== All EVLA antennas have been equipped with L and C band receivers. As of January 2012, 20 EVLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Sensitivity Observational status summary page]. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the EVLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider #what angular resolution is required for your science at the desired observing frequency #for resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission #how much flux will be resolved out by the array configuration that gives the desired angular resolution #will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * To assess if the calibrator is appropriate for the selected bandwidth, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq 4 </math> within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * To assess if the bandpass calibrator is appropriate for the selected channel width, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the the S/N is <math> \geq 4 </math> within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 18d1b6de43e3b855f88919117bb3aa6139a74bc2 881 880 2012-01-25T21:20:28Z Dshepher 15 wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning EVLA observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the EVLA, consult the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for EVLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst EVLA Proposal Submission Tool (PST)]. ===Correlator Setup=== Starting with the D-configuration in 2011 OSRO modes can use up to 8 contiguous subbands per intermediate frequency passband ("baseband") with a maximum (minimum) bandwidth of 128 MHz (0.03 MHz) per subband. All subbands must have the same width and number of channels. There are 2 such independently tunable basebands available. See the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page] for a complete description. You should pick the correlator mode that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see for example [http://en.wikipedia.org/wiki/Gibbs_phenomenon]). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Observing Frequencies=== All EVLA antennas have been equipped with L and C band receivers. As of January 2012, 20 EVLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Sensitivity Observational status summary page]. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the EVLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider #what angular resolution is required for your science at the desired observing frequency #for resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission #how much flux will be resolved out by the array configuration that gives the desired angular resolution #will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * To assess if the calibrator is appropriate for the selected bandwidth, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq 4 </math> within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * To assess if the bandpass calibrator is appropriate for the selected channel width, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the the S/N is <math> \geq 4 </math> within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 74a548e3176a94b51bf2eb6bb2ab3f3d9dfb495a 883 881 2012-01-25T21:33:16Z Dshepher 15 wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning EVLA observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the EVLA, consult the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for EVLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst EVLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All EVLA antennas have been equipped with L and C band receivers. As of January 2012, 20 EVLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Sensitivity Observational status summary page]. ===Correlator Setup=== Starting with the D-configuration in 2011 OSRO modes can use up to 8 contiguous subbands per intermediate frequency passband ("baseband") with a maximum (minimum) bandwidth of 128 MHz (0.03 MHz) per subband. All subbands must have the same width and number of channels. There are 2 such independently tunable basebands available. See the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page] for a complete description. You should pick the correlator mode that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see for example [http://en.wikipedia.org/wiki/Gibbs_phenomenon]). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the EVLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider #what angular resolution is required for your science at the desired observing frequency #for resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission #how much flux will be resolved out by the array configuration that gives the desired angular resolution #will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * To assess if the calibrator is appropriate for the selected bandwidth, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq 4 </math> within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * To assess if the bandpass calibrator is appropriate for the selected channel width, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the the S/N is <math> \geq 4 </math> within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 4eba77ea964d48795589cb3588a0379b98769518 887 883 2012-01-25T21:40:28Z Dshepher 15 /* Correlator Setup */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning EVLA observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the EVLA, consult the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for EVLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst EVLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All EVLA antennas have been equipped with L and C band receivers. As of January 2012, 20 EVLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Sensitivity Observational status summary page]. ===Correlator Setup=== Starting with the D-configuration in 2011 Open Shared Risk Observing (OSRO) modes can use up to 8 contiguous subbands per intermediate frequency passband ("baseband") with a maximum (minimum) bandwidth of 128 MHz (0.03 MHz) per subband. All subbands must have the same width and number of channels. There are 2 such independently tunable basebands available. See the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page] for a complete description. You should pick the correlator mode that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see for example [http://en.wikipedia.org/wiki/Gibbs_phenomenon]). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the EVLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider #what angular resolution is required for your science at the desired observing frequency #for resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission #how much flux will be resolved out by the array configuration that gives the desired angular resolution #will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * To assess if the calibrator is appropriate for the selected bandwidth, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq 4 </math> within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * To assess if the bandpass calibrator is appropriate for the selected channel width, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the the S/N is <math> \geq 4 </math> within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] c1da1ae98eb1aeb6c3a658e61607a4d6a0288b8a 888 887 2012-01-25T21:45:27Z Dshepher 15 /* Correlator Setup */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning EVLA observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the EVLA, consult the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for EVLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst EVLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All EVLA antennas have been equipped with L and C band receivers. As of January 2012, 20 EVLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Sensitivity Observational status summary page]. ===Correlator Setup=== As of September 2011, the Open Shared Risk Observing (OSRO) program offers up to 2 GHz, configured as two independently tunable basebands, each with up to 8 contiguous sub-band pairs of identical bandwidth and channelization. See the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page] for a complete description. You should pick the correlator mode that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see for example [http://en.wikipedia.org/wiki/Gibbs_phenomenon]). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the EVLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider #what angular resolution is required for your science at the desired observing frequency #for resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission #how much flux will be resolved out by the array configuration that gives the desired angular resolution #will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * To assess if the calibrator is appropriate for the selected bandwidth, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq 4 </math> within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * To assess if the bandpass calibrator is appropriate for the selected channel width, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the the S/N is <math> \geq 4 </math> within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] aaf444f1251c8e519cdd19b364f8fa073d808e46 890 888 2012-01-25T22:03:06Z Dshepher 15 /* Correlator Setup */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning EVLA observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the EVLA, consult the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for EVLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst EVLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All EVLA antennas have been equipped with L and C band receivers. As of January 2012, 20 EVLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Sensitivity Observational status summary page]. ===Correlator Setup=== As of September 2011, the Open Shared Risk Observing (OSRO) program offers up to 2 GHz, configured as two independently tunable basebands, each with up to 8 contiguous sub-band pairs of identical bandwidth and channelization. See the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page] for a complete description. You should pick the correlator mode that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the EVLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider #what angular resolution is required for your science at the desired observing frequency #for resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission #how much flux will be resolved out by the array configuration that gives the desired angular resolution #will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * To assess if the calibrator is appropriate for the selected bandwidth, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq 4 </math> within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * To assess if the bandpass calibrator is appropriate for the selected channel width, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the the S/N is <math> \geq 4 </math> within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] cfc148ec5570197a80e1e49fd2a91a7c9a28871b 894 890 2012-01-25T22:17:24Z Dshepher 15 /* Configuration */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning EVLA observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the EVLA, consult the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for EVLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst EVLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All EVLA antennas have been equipped with L and C band receivers. As of January 2012, 20 EVLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Sensitivity Observational status summary page]. ===Correlator Setup=== As of September 2011, the Open Shared Risk Observing (OSRO) program offers up to 2 GHz, configured as two independently tunable basebands, each with up to 8 contiguous sub-band pairs of identical bandwidth and channelization. See the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page] for a complete description. You should pick the correlator mode that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the EVLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * To assess if the calibrator is appropriate for the selected bandwidth, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq 4 </math> within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * To assess if the bandpass calibrator is appropriate for the selected channel width, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the the S/N is <math> \geq 4 </math> within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] db2d1ff44248159f5d7a33c9bb467e50b805fce8 899 894 2012-01-25T22:30:53Z Dshepher 15 /* Estimating Sensitivities */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning EVLA observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the EVLA, consult the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for EVLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst EVLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All EVLA antennas have been equipped with L and C band receivers. As of January 2012, 20 EVLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Sensitivity Observational status summary page]. ===Correlator Setup=== As of September 2011, the Open Shared Risk Observing (OSRO) program offers up to 2 GHz, configured as two independently tunable basebands, each with up to 8 contiguous sub-band pairs of identical bandwidth and channelization. See the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page] for a complete description. You should pick the correlator mode that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the EVLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * To assess if the bandpass calibrator is appropriate for the selected channel width, a good rule of thumb is: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the the S/N is <math> \geq 4 </math> within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 17ed7cc1d9d470ecb89871438474ef3007a11d85 901 899 2012-01-25T22:36:17Z Dshepher 15 /* Estimating Sensitivities */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning EVLA observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the EVLA, consult the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for EVLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst EVLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All EVLA antennas have been equipped with L and C band receivers. As of January 2012, 20 EVLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Sensitivity Observational status summary page]. ===Correlator Setup=== As of September 2011, the Open Shared Risk Observing (OSRO) program offers up to 2 GHz, configured as two independently tunable basebands, each with up to 8 contiguous sub-band pairs of identical bandwidth and channelization. See the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page] for a complete description. You should pick the correlator mode that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the EVLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 3a34971b08289e8607ce86f9a0cf2c9d5acc71c8 904 901 2012-01-25T23:23:23Z Dshepher 15 /* Bandpass Calibration */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning EVLA observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the EVLA, consult the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for EVLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst EVLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All EVLA antennas have been equipped with L and C band receivers. As of January 2012, 20 EVLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Sensitivity Observational status summary page]. ===Correlator Setup=== As of September 2011, the Open Shared Risk Observing (OSRO) program offers up to 2 GHz, configured as two independently tunable basebands, each with up to 8 contiguous sub-band pairs of identical bandwidth and channelization. See the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page] for a complete description. You should pick the correlator mode that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the EVLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 410339eabae023c1d28646b22aa633821f37a0bb 0 48 941 614 2012-02-29T23:27:18Z Jott 8 wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [[ http://splatalogue.net Splatalogue]] which contains data from the [[http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog]], the [[http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [[http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy]] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Line Frequency == The first step is to determine the observing frequency of the spectral line. Two equations may come handy: <math> \nu = \nu_0 / (z+1) </math> determines the (rough) observing frequency <math>\nu</math> as a function of the redshift <math>z</math> and the rest frequency of the spectral line <maths>\nu_0</maths> == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 3397668b2dc2f02fbadde4de5e429967cf3e02dc 942 941 2012-02-29T23:44:54Z Jott 8 wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [[ http://splatalogue.net Splatalogue]] which contains data from the [[http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog]], the [[http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [[http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy]] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Line Frequency == The first step is to determine the observing frequency of the spectral line. Two equations may come handy: <math> \nu = \nu_0 / (z+1) </math> determines the (rough) observing frequency <math>\nu</math> as a function of the redshift <math>z</math> and the rest frequency of the spectral line <maths>\nu_0</maths> == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === df6a6dd0c9ab279d4cde1183054293ec2a13c5d5 943 942 2012-02-29T23:45:28Z Jott 8 /* Observation Planning */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Line Frequency == The first step is to determine the observing frequency of the spectral line. Two equations may come handy: <math> \nu = \nu_0 / (z+1) </math> determines the (rough) observing frequency <math>\nu</math> as a function of the redshift <math>z</math> and the rest frequency of the spectral line <maths>\nu_0</maths> == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === d577d54d52b4e95b07039f21d2d2d53d3bd34352 944 943 2012-03-01T00:13:30Z Jott 8 /* Line Frequency */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Line Frequency == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = (\nu_0^{2} - \nu^{2})/(\nu_0^{2}+\nu^{2})</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = (\lambda_0-\lambda)/(\lambda_0) c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = (\nu-\nu_0)/(\nu_0) c = (\lambda_0-\lambda)/(\lambda) c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \nu_0 / (z+1) </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceeding ones: Correct for: Amplitude (km s$^{-1}$) Rest frame: Nothing added 0.0 topocentric Earth's rotation $\leq$0.5 Earth's motion around $\leq$0.013 km s$^{-1}$ geocentric earth/moon barycenter Earth's motion around $\leq$30 km s$^{-1}$ heliocentric(z) the Sun Solar motion around the $\leq$0.012 km s$^{-1}$ barycentric Solar System barycenter Solar motion $\sim$20 km s$^{-1}$ local standard of rest (LSR) Galactic rotation $\sim$300 km s$^{-1}$ galactocentric The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. More often than not, one wishes to specify the velocity of the object and let the on-line system do the conversion to sky frequency for each scan. This is called ``Doppler tracking''. The user will have to specify, in observe, the rest frequency, choice of rest frame, and radial velocity. Doppler tracking is not implemented during a scan as the frequencies are set at the beginning of each scan. If very accurate tracking is required, one is advised to use short scans. Note that the on-line system uses the same algorithm as dopset whereas observe (version 3 and higher) uses a slightly different method and calculates the observing frequency to within a few tens of meters per sec to the values derived using dopset. Ideally, one wants calibrators observed at the same sky frequency as the sources. This can be achieved by specifying ``no change'' in observe for the flukesynthesizer on the calibrators instead of a velocity. The effect of this is that the LO settings are not changed from what they were during the previous scan. If one wants to start a sequence with a calibrator, it is necessary to precede it with a dummy 1-minute source scan to force the on-line computers to set the LO chain to the proper values. WARNING: if the system crashes and comes back up in the middle of a calibrator scan this scan will be useless because the frequency setting will be in error. The VLA Operator should be alerted to the use of the ``no change'' option, so he can restart an interrupted observation by including again a 1-minute dummy scan. Because of this overhead, it is advised that the use of ``no change'' be restricted to those cases in which it is essential, such as for obtaining high spectral dynamic range; in general, the frequency difference for nearby calibrators is negligable. == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === b5dd2ec16663cff57eaeafffc0f6bbf4233108b2 945 944 2012-03-01T00:14:36Z Jott 8 /* Line Frequency */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Line Frequency == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = (\nu_0^{2} - \nu^{2})/(\nu_0^{2}+\nu^{2})</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0} c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = (\nu-\nu_0)/(\nu_0) c = (\lambda_0-\lambda)/(\lambda) c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \nu_0 / (z+1) </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceeding ones: Correct for: Amplitude (km s$^{-1}$) Rest frame: Nothing added 0.0 topocentric Earth's rotation $\leq$0.5 Earth's motion around $\leq$0.013 km s$^{-1}$ geocentric earth/moon barycenter Earth's motion around $\leq$30 km s$^{-1}$ heliocentric(z) the Sun Solar motion around the $\leq$0.012 km s$^{-1}$ barycentric Solar System barycenter Solar motion $\sim$20 km s$^{-1}$ local standard of rest (LSR) Galactic rotation $\sim$300 km s$^{-1}$ galactocentric The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. More often than not, one wishes to specify the velocity of the object and let the on-line system do the conversion to sky frequency for each scan. This is called ``Doppler tracking''. The user will have to specify, in observe, the rest frequency, choice of rest frame, and radial velocity. Doppler tracking is not implemented during a scan as the frequencies are set at the beginning of each scan. If very accurate tracking is required, one is advised to use short scans. Note that the on-line system uses the same algorithm as dopset whereas observe (version 3 and higher) uses a slightly different method and calculates the observing frequency to within a few tens of meters per sec to the values derived using dopset. Ideally, one wants calibrators observed at the same sky frequency as the sources. This can be achieved by specifying ``no change'' in observe for the flukesynthesizer on the calibrators instead of a velocity. The effect of this is that the LO settings are not changed from what they were during the previous scan. If one wants to start a sequence with a calibrator, it is necessary to precede it with a dummy 1-minute source scan to force the on-line computers to set the LO chain to the proper values. WARNING: if the system crashes and comes back up in the middle of a calibrator scan this scan will be useless because the frequency setting will be in error. The VLA Operator should be alerted to the use of the ``no change'' option, so he can restart an interrupted observation by including again a 1-minute dummy scan. Because of this overhead, it is advised that the use of ``no change'' be restricted to those cases in which it is essential, such as for obtaining high spectral dynamic range; in general, the frequency difference for nearby calibrators is negligable. == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 694cdd5d775104367131e8e7141aae40444c53d9 946 945 2012-03-01T00:16:20Z Jott 8 /* Line Frequency */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Line Frequency == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \nu_0 / (z+1) </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceeding ones: Correct for: Amplitude (km s$^{-1}$) Rest frame: Nothing added 0.0 topocentric Earth's rotation $\leq$0.5 Earth's motion around $\leq$0.013 km s$^{-1}$ geocentric earth/moon barycenter Earth's motion around $\leq$30 km s$^{-1}$ heliocentric(z) the Sun Solar motion around the $\leq$0.012 km s$^{-1}$ barycentric Solar System barycenter Solar motion $\sim$20 km s$^{-1}$ local standard of rest (LSR) Galactic rotation $\sim$300 km s$^{-1}$ galactocentric The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. More often than not, one wishes to specify the velocity of the object and let the on-line system do the conversion to sky frequency for each scan. This is called ``Doppler tracking''. The user will have to specify, in observe, the rest frequency, choice of rest frame, and radial velocity. Doppler tracking is not implemented during a scan as the frequencies are set at the beginning of each scan. If very accurate tracking is required, one is advised to use short scans. Note that the on-line system uses the same algorithm as dopset whereas observe (version 3 and higher) uses a slightly different method and calculates the observing frequency to within a few tens of meters per sec to the values derived using dopset. Ideally, one wants calibrators observed at the same sky frequency as the sources. This can be achieved by specifying ``no change'' in observe for the flukesynthesizer on the calibrators instead of a velocity. The effect of this is that the LO settings are not changed from what they were during the previous scan. If one wants to start a sequence with a calibrator, it is necessary to precede it with a dummy 1-minute source scan to force the on-line computers to set the LO chain to the proper values. WARNING: if the system crashes and comes back up in the middle of a calibrator scan this scan will be useless because the frequency setting will be in error. The VLA Operator should be alerted to the use of the ``no change'' option, so he can restart an interrupted observation by including again a 1-minute dummy scan. Because of this overhead, it is advised that the use of ``no change'' be restricted to those cases in which it is essential, such as for obtaining high spectral dynamic range; in general, the frequency difference for nearby calibrators is negligable. == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 07227f9c081e2b117e9846f4632a4fdc231b8860 947 946 2012-03-01T00:19:04Z Jott 8 /* Line Frequency */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Line Frequency == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceeding ones: Correct for: Amplitude (km s$^{-1}$) Rest frame: Nothing added 0.0 topocentric Earth's rotation $\leq$0.5 Earth's motion around $\leq$0.013 km s$^{-1}$ geocentric earth/moon barycenter Earth's motion around $\leq$30 km s$^{-1}$ heliocentric(z) the Sun Solar motion around the $\leq$0.012 km s$^{-1}$ barycentric Solar System barycenter Solar motion $\sim$20 km s$^{-1}$ local standard of rest (LSR) Galactic rotation $\sim$300 km s$^{-1}$ galactocentric The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. More often than not, one wishes to specify the velocity of the object and let the on-line system do the conversion to sky frequency for each scan. This is called ``Doppler tracking''. The user will have to specify, in observe, the rest frequency, choice of rest frame, and radial velocity. Doppler tracking is not implemented during a scan as the frequencies are set at the beginning of each scan. If very accurate tracking is required, one is advised to use short scans. Note that the on-line system uses the same algorithm as dopset whereas observe (version 3 and higher) uses a slightly different method and calculates the observing frequency to within a few tens of meters per sec to the values derived using dopset. Ideally, one wants calibrators observed at the same sky frequency as the sources. This can be achieved by specifying ``no change'' in observe for the flukesynthesizer on the calibrators instead of a velocity. The effect of this is that the LO settings are not changed from what they were during the previous scan. If one wants to start a sequence with a calibrator, it is necessary to precede it with a dummy 1-minute source scan to force the on-line computers to set the LO chain to the proper values. WARNING: if the system crashes and comes back up in the middle of a calibrator scan this scan will be useless because the frequency setting will be in error. The VLA Operator should be alerted to the use of the ``no change'' option, so he can restart an interrupted observation by including again a 1-minute dummy scan. Because of this overhead, it is advised that the use of ``no change'' be restricted to those cases in which it is essential, such as for obtaining high spectral dynamic range; in general, the frequency difference for nearby calibrators is negligable. == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 7f95d5f3e4c9f077359e90395054c9d3cad3337f 948 947 2012-03-01T00:23:31Z Jott 8 /* Velocity Frames */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Line Frequency == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceeding ones: <tb> <td>Correct for:</td><td> Amplitude (km s$^{-1}$)</td><td> Rest frame: </td></tr> Frame of the Telescope location (nothing added) </td><td> 0.0 </td><td> topocentric</td></tr> Earth's rotation $\leq$0.5 Earth's motion around $\leq$0.013 km s$^{-1}$ geocentric earth/moon barycenter Earth's motion around $\leq$30 km s$^{-1}$ heliocentric(z) the Sun Solar motion around the $\leq$0.012 km s$^{-1}$ barycentric Solar System barycenter Solar motion $\sim$20 km s$^{-1}$ local standard of rest (LSR) Galactic rotation $\sim$300 km s$^{-1}$ galactocentric The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. More often than not, one wishes to specify the velocity of the object and let the on-line system do the conversion to sky frequency for each scan. This is called ``Doppler tracking''. The user will have to specify, in observe, the rest frequency, choice of rest frame, and radial velocity. Doppler tracking is not implemented during a scan as the frequencies are set at the beginning of each scan. If very accurate tracking is required, one is advised to use short scans. Note that the on-line system uses the same algorithm as dopset whereas observe (version 3 and higher) uses a slightly different method and calculates the observing frequency to within a few tens of meters per sec to the values derived using dopset. Ideally, one wants calibrators observed at the same sky frequency as the sources. This can be achieved by specifying ``no change'' in observe for the flukesynthesizer on the calibrators instead of a velocity. The effect of this is that the LO settings are not changed from what they were during the previous scan. If one wants to start a sequence with a calibrator, it is necessary to precede it with a dummy 1-minute source scan to force the on-line computers to set the LO chain to the proper values. WARNING: if the system crashes and comes back up in the middle of a calibrator scan this scan will be useless because the frequency setting will be in error. The VLA Operator should be alerted to the use of the ``no change'' option, so he can restart an interrupted observation by including again a 1-minute dummy scan. Because of this overhead, it is advised that the use of ``no change'' be restricted to those cases in which it is essential, such as for obtaining high spectral dynamic range; in general, the frequency difference for nearby calibrators is negligable. == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 3fb5f327e2b617dba34a9cbea628e90fd85debb2 949 948 2012-03-01T00:39:00Z Jott 8 /* Velocity Frames */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Line Frequency == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceeding ones: <TABLE CELLPADDING=3 BORDER="0"> <td>Rest Frame</td>Correct for:</td><td> Amplitude (km s$^{-1}$)</td><td> Rest frame: </td></tr> nothing added</td><td> 0.0 </td><td> topocentric</td></tr> <td>Earth's rotation $\leq$0.5 <td>Earth's motion around $\leq$0.013 km s$^{-1}$ geocentric earth/moon barycenter <td>Earth's motion around $\leq$30 km s$^{-1}$ heliocentric(z) the Sun <td>Solar motion around the $\leq$0.012 km s$^{-1}$ barycentric Solar System barycenter <td>Solar motion $\sim$20 km s$^{-1}$ local standard of rest (LSR) <td>Galactic rotation $\sim$300 km s$^{-1}$ galactocentric The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. More often than not, one wishes to specify the velocity of the object and let the on-line system do the conversion to sky frequency for each scan. This is called ``Doppler tracking''. The user will have to specify, in observe, the rest frequency, choice of rest frame, and radial velocity. Doppler tracking is not implemented during a scan as the frequencies are set at the beginning of each scan. If very accurate tracking is required, one is advised to use short scans. Note that the on-line system uses the same algorithm as dopset whereas observe (version 3 and higher) uses a slightly different method and calculates the observing frequency to within a few tens of meters per sec to the values derived using dopset. Ideally, one wants calibrators observed at the same sky frequency as the sources. This can be achieved by specifying ``no change'' in observe for the flukesynthesizer on the calibrators instead of a velocity. The effect of this is that the LO settings are not changed from what they were during the previous scan. If one wants to start a sequence with a calibrator, it is necessary to precede it with a dummy 1-minute source scan to force the on-line computers to set the LO chain to the proper values. WARNING: if the system crashes and comes back up in the middle of a calibrator scan this scan will be useless because the frequency setting will be in error. The VLA Operator should be alerted to the use of the ``no change'' option, so he can restart an interrupted observation by including again a 1-minute dummy scan. Because of this overhead, it is advised that the use of ``no change'' be restricted to those cases in which it is essential, such as for obtaining high spectral dynamic range; in general, the frequency difference for nearby calibrators is negligable. == Gibbs Phenomenon and Hanning smoothing == == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S<sub>cal</sub> observed for a time t<sub>cal</sub> and a science target with flux density S<sub>obj</sub> observed for a time t<sub>obj</sub>, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === a4052d5ca1b6c5e0c5d7c86f6a7a4c33517633db 950 949 2012-03-01T00:41:26Z Jott 8 /* Line Frequency */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceeding ones: <TABLE CELLPADDING=3 BORDER="0"> <td>Rest Frame</td>Correct for:</td><td> Amplitude (km s$^{-1}$)</td><td> Rest frame: </td></tr> nothing added</td><td> 0.0 </td><td> topocentric</td></tr> <td>Earth's rotation $\leq$0.5 <td>Earth's motion around $\leq$0.013 km s$^{-1}$ geocentric earth/moon barycenter <td>Earth's motion around $\leq$30 km s$^{-1}$ heliocentric(z) the Sun <td>Solar motion around the $\leq$0.012 km s$^{-1}$ barycentric Solar System barycenter <td>Solar motion $\sim$20 km s$^{-1}$ local standard of rest (LSR) <td>Galactic rotation $\sim$300 km s$^{-1}$ galactocentric The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. More often than not, one wishes to specify the velocity of the object and let the on-line system do the conversion to sky frequency for each scan. This is called ``Doppler tracking''. The user will have to specify, in observe, the rest frequency, choice of rest frame, and radial velocity. Doppler tracking is not implemented during a scan as the frequencies are set at the beginning of each scan. If very accurate tracking is required, one is advised to use short scans. Note that the on-line system uses the same algorithm as dopset whereas observe (version 3 and higher) uses a slightly different method and calculates the observing frequency to within a few tens of meters per sec to the values derived using dopset. Ideally, one wants calibrators observed at the same sky frequency as the sources. This can be achieved by specifying ``no change'' in observe for the flukesynthesizer on the calibrators instead of a velocity. The effect of this is that the LO settings are not changed from what they were during the previous scan. If one wants to start a sequence with a calibrator, it is necessary to precede it with a dummy 1-minute source scan to force the on-line computers to set the LO chain to the proper values. WARNING: if the system crashes and comes back up in the middle of a calibrator scan this scan will be useless because the frequency setting will be in error. The VLA Operator should be alerted to the use of the ``no change'' option, so he can restart an interrupted observation by including again a 1-minute dummy scan. Because of this overhead, it is advised that the use of ``no change'' be restricted to those cases in which it is essential, such as for obtaining high spectral dynamic range; in general, the frequency difference for nearby calibrators is negligable. == Gibbs Phenomenon and Hanning smoothing == == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 4a305400b8a1c2dddfddb7a142aef2c4a866dc89 951 950 2012-03-01T01:05:04Z Jott 8 /* Velocity Frames */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceeding ones: <TABLE CELLPADDING=3 BORDER="0"> <tr><td>Rest Frame Name</td><td>Rest Frame</td><td>Correct for</td><td>Max amplitude [km/s]</td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Gibbs Phenomenon and Hanning smoothing == == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 1530366de1650edf8c3ef83904f63a1cbdd14508 952 951 2012-03-01T01:07:20Z Jott 8 /* Velocity Frames */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceeding ones: <TABLE CELLPADDING=3 BORDER="0"> <tr><th>Rest Frame Name</th><th>Rest Frame</th><th>Correct for</th><th>Max amplitude [km/s]</th></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == == Gibbs Phenomenon and Hanning smoothing == == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === f34a6ff5bb4eda2e0a4f53c21a1d592341826b16 953 952 2012-03-01T01:09:28Z Jott 8 /* Velocity Frames */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="0"> <tr><td><b>Rest Frame Name</td><td>Rest Frame</td><td>Correct for</td><td>Max amplitude [km/s]</td></b></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == == Gibbs Phenomenon and Hanning smoothing == == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === d27398fff76ba7a45ecf764cb854e9c0ce941e24 954 953 2012-03-01T01:10:53Z Jott 8 /* Velocity Frames */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == == Gibbs Phenomenon and Hanning smoothing == == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 8be1dc1d70d1852f650f6901a01c2251aa9bb815 955 954 2012-03-01T01:12:21Z Jott 8 /* Velocity Frames */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == == Gibbs Phenomenon and Hanning smoothing == == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 46144ccce453a546e101fe6dd95555aea506ddbe 956 955 2012-03-01T01:12:56Z Jott 8 /* Velocity Frames */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == == Gibbs Phenomenon and Hanning smoothing == == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 2ea4d796e4f7eea7aedd32d339611f2e2388902c 957 956 2012-03-01T01:13:40Z Jott 8 /* Velocity Frames */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == == Gibbs Phenomenon and Hanning smoothing == == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === f4749e5076fd5b65157af997262aa74255946341 958 957 2012-03-01T01:34:47Z Jott 8 /* Doppler Correction */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tacking''' and was standard for the VLA. the EVLA does NOT support Doppler tracking. The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === e7f89c8d0ade75c1cebded46e964fc08823614a3 959 958 2012-03-01T01:35:19Z Jott 8 /* Doppler Correction */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. the EVLA does NOT support Doppler tracking. The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 9df78942ae3a127e3562d9a900d890db157b0196 960 959 2012-03-01T01:36:10Z Jott 8 /* Gibbs Phenomenon and Hanning smoothing */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. the EVLA does NOT support Doppler tracking. The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == == Sensitivty Calculation == == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === cad57ee979ff996bc8405054c5110ac0415d23b3 961 960 2012-03-01T01:47:34Z Jott 8 /* Gibbs Phenomenon and Hanning smoothing */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. the EVLA does NOT support Doppler tracking. The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivty Calculation == == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 46765d2338c2cd454b5f5eee0262ad0509a292c4 962 961 2012-03-01T01:48:55Z Jott 8 /* Doppler Correction */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivty Calculation == == Setting up Correlator Modes== The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 9c24846fde3551d1dc0215f24dcc4e7d3ff48726 963 962 2012-03-01T02:14:38Z Jott 8 /* Setting up Correlator Modes */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivty Calculation == == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</i>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === db48c4899dcb43318cfcce9ae6f8c0c845710c87 964 963 2012-03-01T02:16:31Z Jott 8 /* Correlator Setup */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivty Calculation == == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. = Basebands = Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</i>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> = Fundamental Subbands and 128 MHz suckouts = = Narrow Subbands = The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 609bb833e14d60625d5adc0dd8df57e120db03ab 965 964 2012-03-01T02:17:11Z Jott 8 /* Correlator Setup */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivty Calculation == = Correlator Setup = The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. = Basebands = Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</i>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> = Fundamental Subbands and 128 MHz suckouts = = Narrow Subbands = The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 70e45f3b641959cce033412af6c8f76e76fb7d2a 966 965 2012-03-01T02:18:19Z Jott 8 /* Correlator Setup */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivty Calculation == === Correlator Setup === The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. = Basebands = Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</i>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> = Fundamental Subbands and 128 MHz suckouts = = Narrow Subbands = The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 7810ca74043b22c12f33f8974f9be245dc835524 967 966 2012-03-01T02:18:39Z Jott 8 /* Correlator Setup */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivty Calculation == == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. = Basebands = Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</i>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> = Fundamental Subbands and 128 MHz suckouts = = Narrow Subbands = The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 2599415d76d8719070fb9dac704719e046f18fea 968 967 2012-03-01T02:19:01Z Jott 8 /* Basebands */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivty Calculation == == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</i>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> = Fundamental Subbands and 128 MHz suckouts = = Narrow Subbands = The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 44b675feac93b6e928601b7da884a316fa1e719d 969 968 2012-03-01T02:19:26Z Jott 8 /* Fundamental Subbands and 128 MHz suckouts */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivty Calculation == == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</i>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> === Fundamental Subbands and 128 MHz suckouts === = Narrow Subbands = The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 9387dcdbbb08f8f23bc5207f01a516d39fa3f058 970 969 2012-03-01T02:19:56Z Jott 8 /* Basebands */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivty Calculation == == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> === Fundamental Subbands and 128 MHz suckouts === = Narrow Subbands = The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === f7be4e4bbf5dfe9bdb62ae7e0d6f300817f65d92 971 970 2012-03-01T02:20:23Z Jott 8 /* Narrow Subbands */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivty Calculation == == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> === Fundamental Subbands and 128 MHz suckouts === === Narrow Subbands === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 2bd073363e23b0563d3f8aa6fdea7b28809f6d9d 972 971 2012-03-01T02:24:14Z Jott 8 /* Narrow Subbands */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivty Calculation == == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> === Fundamental Subbands and 128 MHz suckouts === === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. They can be any <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </tbody> </table> <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </tbody> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 9b562149726b3d148d0a29bce508b55199cc2362 973 972 2012-03-01T02:29:23Z Jott 8 /* Narrow Subbands */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivty Calculation == == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> === Fundamental Subbands and 128 MHz suckouts === === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </tbody> </table> * Dual Polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </tbody> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 9c9a16dff9ba084e10d1556b0f12c83e144343fe 974 973 2012-03-06T00:03:29Z Jott 8 /* Sensitivty Calculation */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the '''low frequency guide''' (??? insert link ???) and '''high frequency guide''' (???? insert link ???) for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> === Fundamental Subbands and 128 MHz suckouts === === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </tbody> </table> * Dual Polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </tbody> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === f4663d47d087d5ff436cad4b44caf2fce183b418 0 48 975 974 2012-03-06T00:07:27Z Jott 8 /* Fundamental Subbands and 128 MHz suckouts */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the '''low frequency guide''' (??? insert link ???) and '''high frequency guide''' (???? insert link ???) for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> === Fundamental Subbands and 128 MHz suckouts === After the analog filter of the basebands, the signal enters the correlator and 128MHz wide fundamental subbands. They are placed adjacently to cover the full width of the basebands (8 fundamental subbands per 1GHz baseband, see figure ???). As each fundamental subband has some filter shape with soft corners, this implies that the sensitivity of the EVLA drops to about half its value between any two fundamental subbands. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </tbody> </table> * Dual Polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </tbody> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === d0ace3fa63e8bce85ce13eee10688654b7159940 976 975 2012-03-06T00:21:43Z Jott 8 wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the '''low frequency guide''' (??? insert link ???) and '''high frequency guide''' (???? insert link ???) for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> === Fundamental Subbands and 128 MHz suckouts === After the analog filter of the basebands, the signal enters the correlator and 128MHz wide fundamental subbands. They are placed adjacently to cover the full width of the basebands (8 fundamental subbands per 1GHz baseband, see figure ???). As each fundamental subband has some filter shape with soft corners, this implies that the sensitivity of the EVLA drops to about half its value between any two fundamental subbands. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </tbody> </table> * Dual Polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </tbody> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === a8ac71cce3a6fb0aab0900043d08eb67e293c6c7 977 976 2012-03-06T01:57:43Z Jott 8 /* Doppler Correction */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the '''low frequency guide''' (??? insert link ???) and '''high frequency guide''' (???? insert link ???) for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> === Fundamental Subbands and 128 MHz suckouts === After the analog filter of the basebands, the signal enters the correlator and 128MHz wide fundamental subbands. They are placed adjacently to cover the full width of the basebands (8 fundamental subbands per 1GHz baseband, see figure ???). As each fundamental subband has some filter shape with soft corners, this implies that the sensitivity of the EVLA drops to about half its value between any two fundamental subbands. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </tbody> </table> * Dual Polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </tbody> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === cbd3403c834481fbd1c1a449b2a0194929c3d20a 978 977 2012-03-07T03:17:46Z Jott 8 /* Sensitivity Calculation */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [Category:LowFrequency low frequency guide] and [Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> === Fundamental Subbands and 128 MHz suckouts === After the analog filter of the basebands, the signal enters the correlator and 128MHz wide fundamental subbands. They are placed adjacently to cover the full width of the basebands (8 fundamental subbands per 1GHz baseband, see figure ???). As each fundamental subband has some filter shape with soft corners, this implies that the sensitivity of the EVLA drops to about half its value between any two fundamental subbands. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </tbody> </table> * Dual Polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </tbody> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 6ee59e235a1590c0d2e5b8497dfddfa62f7cc716 979 978 2012-03-07T03:18:27Z Jott 8 /* Sensitivity Calculation */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands with 1 GHz bandwidth each. As a second step, the basebands enter digital filters, 128MHz wide, the fundamental subbands. <img src="https://docs.google.com/drawings/pub?id=1AJjdtrFNRvddyPqyvoJz9EALmBzDlgFPw4-Yu-u2qmc&amp;w=960&amp;h=720"> === Fundamental Subbands and 128 MHz suckouts === After the analog filter of the basebands, the signal enters the correlator and 128MHz wide fundamental subbands. They are placed adjacently to cover the full width of the basebands (8 fundamental subbands per 1GHz baseband, see figure ???). As each fundamental subband has some filter shape with soft corners, this implies that the sensitivity of the EVLA drops to about half its value between any two fundamental subbands. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </tbody> </table> * Dual Polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </tbody> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === c53ae7236d4df51c15a5d24de10b82efcc035c12 981 979 2012-03-07T03:31:36Z Jott 8 /* Basebands */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely inependently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fundamental Subbands and 128 MHz suckouts === After the analog filter of the basebands, the signal enters the correlator and 128MHz wide fundamental subbands. They are placed adjacently to cover the full width of the basebands (8 fundamental subbands per 1GHz baseband, see figure ???). As each fundamental subband has some filter shape with soft corners, this implies that the sensitivity of the EVLA drops to about half its value between any two fundamental subbands. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </tbody> </table> * Dual Polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </tbody> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === d41e24686b93493e0e08440c1356c47dc31defe7 984 981 2012-03-07T03:45:57Z Jott 8 /* Fundamental Subbands and 128 MHz suckouts */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely inependently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </tbody> </table> * Dual Polarization <table class="grid listing"> <tbody> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </tbody> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === f58b6646fbff68ce00de5fc8b4443dc17fb318f4 985 984 2012-03-07T03:46:50Z Jott 8 /* Narrow Subbands */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely inependently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === bf2e31a472e17c5e4aacfb95e70bc2de51b3f334 986 985 2012-03-07T03:47:41Z Jott 8 /* Basebands */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|400pix]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely inependently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === d079f7000e9ee9f0363317f52bfcaff6983c4ba6 987 986 2012-03-07T03:48:00Z Jott 8 /* Basebands */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|200pix]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely inependently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 2036529b5b8902f6a7b3bc2f694f14e4aef0cb34 988 987 2012-03-07T03:49:04Z Jott 8 /* Basebands */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|400px|thumb|left|WIDAR correlator baseband with subbands]]]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely inependently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === a5ad4bee2c720a1798554227094a18e5d48d58f2 989 988 2012-03-07T03:50:05Z Jott 8 /* Basebands */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 621ee18268e7eec66ac4b602a5cd6b5e5a711073 990 989 2012-03-07T03:51:44Z Jott 8 /* Fixed 128MHz Subbands and 128 MHz "Suckouts" */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === f2d3ca5d45386123cf6015d46d5b0a20a299315c 991 990 2012-03-07T03:53:54Z Jott 8 /* Narrow Subbands */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === fc80dcf46e7d300af2a4a3c3ba81651580a32bb2 992 991 2012-03-07T03:55:14Z Jott 8 /* Narrow Subbands */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 382733673e82989a703c24b04b9106ce5e887a57 993 992 2012-03-07T04:06:18Z Jott 8 /* Baselineboard Stacking */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 8ee241f252f08cb67cbd002e7a25ee23ff2870ee 994 993 2012-03-07T04:09:55Z Jott 8 /* Recirculation */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === The OPT spectral line setup - planning - PST - OPT - exposure calculator - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === f3492b4f97f0ee54669f343861b1c8805b8dfa48 995 994 2012-03-07T04:14:33Z Jott 8 /* Data Rate Limits */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. === Planning and Setup Tools === ==== The Proposal Submission Tool (PST) ==== === Setting up a Spectral Observation in the Obersvation Preparation Tool (OPT) === - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 350faa7557e39bb8d61e486fdf54d9c02e27beb6 996 995 2012-03-07T04:14:57Z Jott 8 /* Setting up a Spectral Observation in the Obersvation Preparation Tool (OPT) */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. === Planning and Setup Tools === ==== The Proposal Submission Tool (PST) ==== ==== Setting up a Spectral Observation in the Obersvation Preparation Tool (OPT) ==== - OPT preparation and setup . suckouts . filters . tuneability of subbands . channels full/dual/single pol . baselineboard stacking . recirculation . gapfree setups . data rate considerations . setups with variable bandwidths - bandpass - post processing - hanning smoothing - velocity systems - continuum subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === dee448d4ddac480cfeb184d3cb1bcaf386c415cd 997 996 2012-03-07T04:16:50Z Jott 8 /* Setting up a Spectral Observation in the Obersvation Preparation Tool (OPT) */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. === Planning and Setup Tools === ==== The Proposal Submission Tool (PST) ==== ==== Setting up a Spectral Observation in the Obersvation Preparation Tool (OPT) ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. * Continuum Subtraction = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === ee47681c11dcd72f9934bae58d6d734ec9dcfea0 998 997 2012-03-07T04:22:50Z Jott 8 /* Planning and Setup Tools */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. === Planning and Setup === ==== The Proposal Submission Tool (PST) ==== ==== Setting up a Spectral Observation using the Observation Preparation Tool (OPT) ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. * Continuum Subtraction * High Dynamic Range imaging For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === d7ae41b25bbdc7b62a50074b19135cb338e7c18e 999 998 2012-03-07T04:23:25Z Jott 8 /* Planning and Setup */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. * Continuum Subtraction * High Dynamic Range imaging For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Detailed Guidelines == === Observing Preparation Recommendations === ==== Scheduling ==== ==== Calibration Strategy ==== * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Monitoring Observations ==== === Post-processing Guidelines === 5598a5fc122ad56eff86eaf6f83e5699e65c3846 1000 999 2012-03-07T04:25:47Z Jott 8 /* Detailed Guidelines */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. * Continuum Subtraction * High Dynamic Range imaging For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. 82b8f28b4ba45fa69cd747d9f47e3a3313d5eeb3 1001 1000 2012-03-07T04:28:23Z Jott 8 /* Velocity Frames */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a veloocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600 </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. * Continuum Subtraction * High Dynamic Range imaging For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. b150162034a283bc18f5629385c1b9b517099f8e 1002 1001 2012-03-07T04:29:05Z Jott 8 /* Velocity Frames */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == Correlator Setup == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. * Continuum Subtraction * High Dynamic Range imaging For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. 1fa0308741a0a8ac220ff7aefe8d08538f7794fc 1003 1002 2012-03-07T04:29:36Z Jott 8 /* Correlator Setup */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. * Continuum Subtraction * High Dynamic Range imaging For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. eb9fa71ec3a698abb5ddb5e99821a22f41fb480a 1004 1003 2012-03-07T04:39:50Z Jott 8 /* Setting up a Spectral Observation using the Observation Preparation Tool (OPT) */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. * Continuum Subtraction The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. * High Dynamic Range imaging For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. d70271540f1602e2d4609226ead103b9d4062c2f 1005 1004 2012-03-07T17:16:03Z Jott 8 /* Setting up a Spectral Observation using the Observation Preparation Tool (OPT) */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. * Continuum Subtraction The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. * High Dynamic Range imaging For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. 3e0e714c448e4399a4dd12895c8b9a8d719777f4 1006 1005 2012-03-07T18:42:22Z Jott 8 /* Tuning Restrictions */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. * Continuum Subtraction The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. * High Dynamic Range imaging For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. b6cc4609fb0c69c13543a53bde326e150b0c6066 1007 1006 2012-03-07T18:43:56Z Jott 8 /* Correlator Baselineboards */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == Observation Planning == The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations. There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. * Continuum Subtraction The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. * High Dynamic Range imaging For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. 7d7a791195a4e2c17be7ac4f95ed5c70fcf76355 1008 1007 2012-03-07T19:29:51Z Jott 8 /* Observation Planning */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects === Line Rest Frequencies === There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. * Bandpass Setup All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. * Continuum Subtraction The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. * High Dynamic Range imaging For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. 95750841b5ebc5f3d4148fab1c2dd3017e8f4a22 1011 1008 2012-03-07T19:37:58Z Jott 8 /* Setting up a Spectral Observation using the Observation Preparation Tool (OPT) */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects === Line Rest Frequencies === There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Bandpass Setup ==== All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. 46a0dd1bd028b60d901f4d771a3b089772ca7a0f 1012 1011 2012-03-07T20:18:42Z Jott 8 /* Setting up a Spectral Observation using the Observation Preparation Tool (OPT) */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects === Line Rest Frequencies === There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== ==== Bandpass Setup ==== All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. 4dece3464b5c114c2c305fbd3e0d385676d1c9bf 1013 1012 2012-03-07T22:55:55Z Jott 8 /* Bandpass Setup */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects === Line Rest Frequencies === There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== ==== Bandpass Setup ==== All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. c458728f4e403253f3f82170b525c1a060091807 1014 1013 2012-03-07T23:10:38Z Jott 8 wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects === Line Rest Frequencies === There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. d0a58551ef324e15a5c3fb27e9b17d2b027b297e 1015 1014 2012-03-07T23:11:54Z Jott 8 /* Doppler Setting */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects === Line Rest Frequencies === There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?]. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. 24957bc8f9752908be076cf551ef9a0f9f6c0abf 1016 1015 2012-03-07T23:13:45Z Jott 8 /* Bandpass Setup */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects === Line Rest Frequencies === There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. 3be7726cf82e6ee74fb91fc894d66795d41b0095 1017 1016 2012-03-07T23:18:17Z Jott 8 /* Bandpass Setup */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects === Line Rest Frequencies === There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. ae13de8609cba5aa8741f46e6f4e170900f17cd1 1018 1017 2012-03-07T23:20:07Z Jott 8 /* Bandpass Setup */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects === Line Rest Frequencies === There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. 3f90b76187270361dea626e75eb080d724e788f2 1019 1018 2012-03-07T23:25:30Z Jott 8 /* Bandpass Setup */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects === Line Rest Frequencies === There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. For that to work, both the narrow and wide subbands need to be at exactly the same center frequency. Different centers ''will'' introduce a phase offset between them. (??? check with Vivek ???) ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. 17fa0bde765c476a69f9d83c6c972ab8683c3491 1020 1019 2012-03-07T23:27:03Z Jott 8 /* Phase/Complex Gain Calibration */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects === Line Rest Frequencies === There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. For that to work, both the narrow and wide subbands need to be at exactly the same center frequency. Different centers ''will'' introduce a phase offset between them. '''(??? check with Vivek ???)''' ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. 9ef58ff3989955df74310ae0f2e26775758ab5fb 1021 1020 2012-03-07T23:29:41Z Jott 8 /* Line Rest Frequencies */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. For that to work, both the narrow and wide subbands need to be at exactly the same center frequency. Different centers ''will'' introduce a phase offset between them. '''(??? check with Vivek ???)''' ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. bbda068ebc07c68ff1847707e1a0db1aca27929f 1022 1021 2012-03-07T23:31:01Z Jott 8 /* High Dynamic Range Imaging */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. For that to work, both the narrow and wide subbands need to be at exactly the same center frequency. Different centers ''will'' introduce a phase offset between them. '''(??? check with Vivek ???)''' ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. 00b0fd559d037f311290568a3c1ea62a1bcce8d6 1023 1022 2012-03-07T23:43:13Z Jott 8 /* Tuning Restrictions */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. For that to work, both the narrow and wide subbands need to be at exactly the same center frequency. Different centers ''will'' introduce a phase offset between them. '''(??? check with Vivek ???)''' ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. 3564bf1071b3411f8f6e1232904151e87245cd6c 1024 1023 2012-03-07T23:43:55Z Jott 8 /* WIDAR Tuning Restrictions */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. For that to work, both the narrow and wide subbands need to be at exactly the same center frequency. Different centers ''will'' introduce a phase offset between them. '''(??? check with Vivek ???)''' ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. ba5060bdd6a25aa8e12b50ffb7156dde381b7bff 6 81 980 2012-03-07T03:23:13Z Jott 8 schematic WIDAR bands wikitext text/x-wiki schematic WIDAR bands 658fecdc5e637c26c0f86b8ca77938d73848860b 6 82 982 2012-03-07T03:44:45Z Jott 8 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 83 983 2012-03-07T03:45:19Z Jott 8 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 84 1009 2012-03-07T19:35:04Z Jott 8 OPT - Instrument Configuration: Baseband settings wikitext text/x-wiki OPT - Instrument Configuration: Baseband settings 1406c9dda2809fd1c83bbc660a1a9dd25794c7b9 6 85 1010 2012-03-07T19:37:09Z Jott 8 OPT - Instrument Configuration: Subband Settings wikitext text/x-wiki OPT - Instrument Configuration: Subband Settings a627b4882aada90a6a325b78e821b36fc511ef0d 0 48 1025 1024 2012-03-08T00:57:18Z Jott 8 /* EVLA Spectral Line Observing */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package the main data reduction software for EVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. For that to work, both the narrow and wide subbands need to be at exactly the same center frequency. Different centers ''will'' introduce a phase offset between them. '''(??? check with Vivek ???)''' ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. f723b5743086c1382f2ddea1122a99c6ecb1d41d 1026 1025 2012-03-08T01:04:49Z Jott 8 /* EVLA Spectral Line Observing */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correrlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package the main data reduction software for EVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage] and the [http://casaguides.nrao.edu CASAguides wiki] contains guides on EVLA spectral line data reduction as well as some hints, tipes & tricks on using CASA and the visulaization tools, such as the casa viewer that are designed to display spectral data cubes == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. For that to work, both the narrow and wide subbands need to be at exactly the same center frequency. Different centers ''will'' introduce a phase offset between them. '''(??? check with Vivek ???)''' ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. a98cdfeaae25e9da096ca9d2b276a2a61f1752b2 1027 1026 2012-03-08T01:06:42Z Jott 8 /* EVLA Spectral Line Observing */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package the main data reduction software for EVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage] and the [http://casaguides.nrao.edu CASAguides wiki] contains guides on EVLA spectral line data reduction as well as some hints, tips & tricks on using CASA and the visualization tools, such as the [http://casa.nrao.edu/CasaViewerDemo/casaViewerDemo.html CASA viewer] that are designed to display spectral data cubes == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. For that to work, both the narrow and wide subbands need to be at exactly the same center frequency. Different centers ''will'' introduce a phase offset between them. '''(??? check with Vivek ???)''' ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. b7bfda3c67f476d97a7261a50edefb11efc962f0 1028 1027 2012-03-08T01:43:02Z Jott 8 /* Phase/Complex Gain Calibration */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package the main data reduction software for EVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage] and the [http://casaguides.nrao.edu CASAguides wiki] contains guides on EVLA spectral line data reduction as well as some hints, tips & tricks on using CASA and the visualization tools, such as the [http://casa.nrao.edu/CasaViewerDemo/casaViewerDemo.html CASA viewer] that are designed to display spectral data cubes == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. For that to work, both the narrow and wide subbands need to be at exactly the same center frequency. In addition, there will be a phase offset between the narrow and wide subbands that needs to be determined and corrected for in post-processing. So unless there's no alternative to this procedure, it is recommended to use exactly the same spectral setup for the complex gain calibrator to what is used for the target sources. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. 5c80bef19c07e23e204995be8d68753474570544 1029 1028 2012-03-08T04:09:37Z Jott 8 /* Phase/Complex Gain Calibration */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package the main data reduction software for EVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage] and the [http://casaguides.nrao.edu CASAguides wiki] contains guides on EVLA spectral line data reduction as well as some hints, tips & tricks on using CASA and the visualization tools, such as the [http://casa.nrao.edu/CasaViewerDemo/casaViewerDemo.html CASA viewer] that are designed to display spectral data cubes == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. dd4abde0b72a0a63e12f3ca70795328b50f201a9 1030 1029 2012-03-08T04:20:22Z Jott 8 /* Bandpass Setup */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package the main data reduction software for EVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage] and the [http://casaguides.nrao.edu CASAguides wiki] contains guides on EVLA spectral line data reduction as well as some hints, tips & tricks on using CASA and the visualization tools, such as the [http://casa.nrao.edu/CasaViewerDemo/casaViewerDemo.html CASA viewer] that are designed to display spectral data cubes == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. d136d38866ed63ccebc8d45ab50f198e11fe303b 1031 1030 2012-03-08T04:21:25Z Jott 8 /* Spectral Line Observing */ wikitext text/x-wiki The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == EVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the EVLA. The new, wide bandwidths of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The EVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the EVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The EVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package the main data reduction software for EVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage] and the [http://casaguides.nrao.edu CASAguides wiki] contains guides on EVLA spectral line data reduction as well as some hints, tips & tricks on using CASA and the visualization tools, such as the [http://casa.nrao.edu/CasaViewerDemo/casaViewerDemo.html CASA viewer] that are designed to display spectral data cubes == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi EVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The EVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the EVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the EVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions EVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the EVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the EVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines EVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and EVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The EVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the EVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the EVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the EVLA (last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive EVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process EVLA spectra line data. dc63be251dea329e82bbbbd2b5a1050fa184fecd 10 2 1032 876 2012-03-29T18:28:19Z Dshepher 15 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary '''Observational Status Summary'''] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:SpectraLine| Spectral Line Observations]] · [[:Category:Polarimetry| Polarimetry Observations]] · [[:Category:Planetary| Planetary Observations]] · [[:Category:HighFrequency| High Frequency Observing (K, Ka, Q)]] · [[:Category:LowFrequency| Low Frequency Observing (L, S, C)]] · [[:Category:PhasedArray| Phased Array Observing]] [[:Category:Pulsar| Pulsar Observing]] · [[:Category:OPT-QuickStart| OPT Quick Start Guide]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> CASA: [http://casaguides.nrao.edu '''CASA Reduction Guides'''] <br> AIPS: [[Key to Calcodes]] · [[:Category:Post-Processing|Special Considerations for EVLA Data Calibration and Imaging in AIPS]]<br> |} c33c9f816a16509d7ee8f4ebcbb122048249bfac 14 86 1033 2012-03-29T18:29:41Z Dshepher 15 Created page with "Place holder for the OPT quick Start guide currently under development by Amanda Kepley at: https://www.evernote.com/shard/s84/sh/a7779bd3-62ca-4da5-bbe1-d5901b71736f/39a8d2bfb..." wikitext text/x-wiki Place holder for the OPT quick Start guide currently under development by Amanda Kepley at: https://www.evernote.com/shard/s84/sh/a7779bd3-62ca-4da5-bbe1-d5901b71736f/39a8d2bfb4ed5573f148b30f3fc308e0 4d954716600ad3c97208add5aeac6f90b0e6e8fc 1034 1033 2012-04-11T23:06:24Z Akepley 16 wikitext text/x-wiki Place holder for the OPT quick Start guide currently under development by Amanda Kepley at: https://www.evernote.com/shard/s84/sh/a7779bd3-62ca-4da5-bbe1-d5901b71736f/39a8d2bfb4ed5573f148b30f3fc308e0 ---- IGNORE THE BELOW----- NEED TO INSERT TOC. ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have e0a2ca56a3eb8cb025cfddfbbd965b95ee81bdd3 1037 1034 2012-04-11T23:27:26Z Akepley 16 wikitext text/x-wiki Place holder for the OPT quick Start guide currently under development by Amanda Kepley at: https://www.evernote.com/shard/s84/sh/a7779bd3-62ca-4da5-bbe1-d5901b71736f/39a8d2bfb4ed5573f148b30f3fc308e0 ---- IGNORE THE BELOW----- NEED TO INSERT TOC. ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. 695aae7c074a3e2508cc0858ee1e50168b619ce9 1038 1037 2012-04-11T23:28:46Z Akepley 16 wikitext text/x-wiki Place holder for the OPT quick Start guide currently under development by Amanda Kepley at: https://www.evernote.com/shard/s84/sh/a7779bd3-62ca-4da5-bbe1-d5901b71736f/39a8d2bfb4ed5573f148b30f3fc308e0 ---- IGNORE THE BELOW----- NEED TO INSERT TOC. ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. 0bed5c83ff7a5e4bf65ed38e9ee8c1cf986f7748 1039 1038 2012-04-11T23:29:27Z Akepley 16 wikitext text/x-wiki Place holder for the OPT quick Start guide currently under development by Amanda Kepley at: https://www.evernote.com/shard/s84/sh/a7779bd3-62ca-4da5-bbe1-d5901b71736f/39a8d2bfb4ed5573f148b30f3fc308e0 ---- IGNORE THE BELOW----- NEED TO INSERT TOC. ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. 7947ff870f4fba2a8bf2fcc235d297276ec33290 1040 1039 2012-04-11T23:30:41Z Akepley 16 wikitext text/x-wiki Place holder for the OPT quick Start guide currently under development by Amanda Kepley at: https://www.evernote.com/shard/s84/sh/a7779bd3-62ca-4da5-bbe1-d5901b71736f/39a8d2bfb4ed5573f148b30f3fc308e0 ---- IGNORE THE BELOW----- NEED TO INSERT TOC. ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]]<br /> The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. fa67246fe490e862b75ecafdf48a1893619cd8cf 1042 1040 2012-04-11T23:33:01Z Akepley 16 /* OPT access and layout */ wikitext text/x-wiki Place holder for the OPT quick Start guide currently under development by Amanda Kepley at: https://www.evernote.com/shard/s84/sh/a7779bd3-62ca-4da5-bbe1-d5901b71736f/39a8d2bfb4ed5573f148b30f3fc308e0 ---- IGNORE THE BELOW----- NEED TO INSERT TOC. ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]]<br /> The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. 7e7db1e59e811c37c25bc539dc99a55bb3e30c7b 1043 1042 2012-04-11T23:39:05Z Akepley 16 wikitext text/x-wiki Place holder for the OPT quick Start guide currently under development by Amanda Kepley at: https://www.evernote.com/shard/s84/sh/a7779bd3-62ca-4da5-bbe1-d5901b71736f/39a8d2bfb4ed5573f148b30f3fc308e0 ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]]<br /> The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. 4b62bc3b58bb25761c535db45369b3bb6cee2657 1044 1043 2012-04-11T23:39:56Z Akepley 16 /* OPT access and layout */ wikitext text/x-wiki Place holder for the OPT quick Start guide currently under development by Amanda Kepley at: https://www.evernote.com/shard/s84/sh/a7779bd3-62ca-4da5-bbe1-d5901b71736f/39a8d2bfb4ed5573f148b30f3fc308e0 ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]]<br /> The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. aefc00a5c1227d8abec4ae7f6c193d6917f885a3 1045 1044 2012-04-12T20:27:38Z Akepley 16 wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]]<br /> The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]]<br /> The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]]<br /> When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]]<br /> All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. * [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. * [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. * [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.√Ǭ†<br /> * [[Image:phase_calibrator_select_group.png]]<br /> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.''√Ǭ†You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel''√Ǭ†in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. √Ǭ†Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the√Ǭ†[http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: ** *** The first choice you have to make is whether or not you are going to use doppler setting. If you're using doppler setting, then select "rest" next to the baseband frequency. If you are not using doppler setting then choose "sky."<br /> *** If doppler setting, **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are√Ǭ†6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 √Ǭ†MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it √Ǭ†and have a format like 12A-186√Ǭ†([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees) for reference pointing to be effective. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]]<br /> The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br />''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page.<br /> To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote and exported to mediawiki. The illustrations were created using Skitch. 6f58af91efca408ae85cf0952e47c2a2369c2836 1046 1045 2012-04-12T20:29:12Z Akepley 16 /* About This Document */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]]<br /> The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]]<br /> The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]]<br /> When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]]<br /> All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. * [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. * [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. * [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.√Ǭ†<br /> * [[Image:phase_calibrator_select_group.png]]<br /> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.''√Ǭ†You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel''√Ǭ†in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. √Ǭ†Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the√Ǭ†[http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: ** *** The first choice you have to make is whether or not you are going to use doppler setting. If you're using doppler setting, then select "rest" next to the baseband frequency. If you are not using doppler setting then choose "sky."<br /> *** If doppler setting, **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are√Ǭ†6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 √Ǭ†MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it √Ǭ†and have a format like 12A-186√Ǭ†([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees) for reference pointing to be effective. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]]<br /> The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br />''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page.<br /> To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. 3da867ea611a55e33d4adfec34a68013f45a4daa 1047 1046 2012-04-12T20:39:19Z Akepley 16 /* Introduction */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]]<br /> The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]]<br /> The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]]<br /> When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]]<br /> All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. * [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. * [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. * [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.√Ǭ†<br /> * [[Image:phase_calibrator_select_group.png]]<br /> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.''√Ǭ†You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel''√Ǭ†in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. √Ǭ†Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the√Ǭ†[http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: ** *** The first choice you have to make is whether or not you are going to use doppler setting. If you're using doppler setting, then select "rest" next to the baseband frequency. If you are not using doppler setting then choose "sky."<br /> *** If doppler setting, **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are√Ǭ†6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 √Ǭ†MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it √Ǭ†and have a format like 12A-186√Ǭ†([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees) for reference pointing to be effective. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]]<br /> The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br />''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page.<br /> To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. 0ce931e502c59dbf9083f55d441d93ba7d7de7ff 1048 1047 2012-04-12T20:41:38Z Akepley 16 /* OPT access and layout */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]]<br /> The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]]<br /> When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]]<br /> All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. * [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. * [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. * [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.√Ǭ†<br /> * [[Image:phase_calibrator_select_group.png]]<br /> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.''√Ǭ†You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel''√Ǭ†in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. √Ǭ†Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the√Ǭ†[http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: ** *** The first choice you have to make is whether or not you are going to use doppler setting. If you're using doppler setting, then select "rest" next to the baseband frequency. If you are not using doppler setting then choose "sky."<br /> *** If doppler setting, **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are√Ǭ†6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 √Ǭ†MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it √Ǭ†and have a format like 12A-186√Ǭ†([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees) for reference pointing to be effective. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]]<br /> The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br />''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page.<br /> To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. 9e700b957d45b6b8df1425745b817b7c9cfe69c5 1049 1048 2012-04-12T20:43:29Z Akepley 16 /* Creating a Source List */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]]<br /> When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]]<br /> All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. * [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. * [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. * [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.√Ǭ†<br /> * [[Image:phase_calibrator_select_group.png]]<br /> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.''√Ǭ†You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel''√Ǭ†in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. √Ǭ†Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the√Ǭ†[http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: ** *** The first choice you have to make is whether or not you are going to use doppler setting. If you're using doppler setting, then select "rest" next to the baseband frequency. If you are not using doppler setting then choose "sky."<br /> *** If doppler setting, **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are√Ǭ†6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 √Ǭ†MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it √Ǭ†and have a format like 12A-186√Ǭ†([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees) for reference pointing to be effective. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]]<br /> The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br />''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page.<br /> To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. b1c50155cbf62a3e38b7559647fd2a39c099c0ae 1050 1049 2012-04-12T20:46:07Z Akepley 16 /* Selecting a Complex Gain Calibrator */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]]<br /> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.''√Ǭ†You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel''√Ǭ†in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. √Ǭ†Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the√Ǭ†[http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: ** *** The first choice you have to make is whether or not you are going to use doppler setting. If you're using doppler setting, then select "rest" next to the baseband frequency. If you are not using doppler setting then choose "sky."<br /> *** If doppler setting, **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are√Ǭ†6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 √Ǭ†MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it √Ǭ†and have a format like 12A-186√Ǭ†([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees) for reference pointing to be effective. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]]<br /> The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br />''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page.<br /> To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. b5f40955af4184fbe6da0215f002a9126f76230b 1051 1050 2012-04-12T20:48:02Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]]<br /> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. √Ǭ†Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the√Ǭ†[http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: ** *** The first choice you have to make is whether or not you are going to use doppler setting. If you're using doppler setting, then select "rest" next to the baseband frequency. If you are not using doppler setting then choose "sky."<br /> *** If doppler setting, **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are√Ǭ†6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 √Ǭ†MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it √Ǭ†and have a format like 12A-186√Ǭ†([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees) for reference pointing to be effective. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]]<br /> The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br />''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page.<br /> To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. d6946eea372e8a217df90ec93fd59d96b2a70e71 1052 1051 2012-04-12T20:50:49Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]]<br /> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the√Ǭ†[http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** If you're using doppler setting, then select "rest" next to the baseband frequency. **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are√Ǭ†6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 √Ǭ†MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it √Ǭ†and have a format like 12A-186√Ǭ†([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees) for reference pointing to be effective. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]]<br /> The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br />''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page.<br /> To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. dbf1d6437f96b6fabce1c51889f787f1bb0aedea 1053 1052 2012-04-12T20:55:17Z Akepley 16 /* Putting It All Together: Creating a Scheduling Block */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]]<br /> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the√Ǭ†[http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** If you're using doppler setting, then select "rest" next to the baseband frequency. **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are√Ǭ†6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 √Ǭ†MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page.<br /> To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. c47d195dfe5d430e3762de9305cb0060c5ef30f8 1054 1053 2012-04-12T20:55:49Z Akepley 16 /* Checking Your Scheduling Block */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]]<br /> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the√Ǭ†[http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** If you're using doppler setting, then select "rest" next to the baseband frequency. **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are√Ǭ†6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 √Ǭ†MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page.<br /> To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. 176dbe1a66dc32afd735d6749b1653aa55208410 1055 1054 2012-04-12T20:56:27Z Akepley 16 /* Submitting Your Scheduling Block */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]]<br /> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the√Ǭ†[http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** If you're using doppler setting, then select "rest" next to the baseband frequency. **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are√Ǭ†6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 √Ǭ†MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. c84301fa3d99f8bb1b9355af1cb6fbf90e31016d 1056 1055 2012-04-12T20:56:55Z Akepley 16 /* Submitting Your Scheduling Block */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]]<br /> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the√Ǭ†[http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** If you're using doppler setting, then select "rest" next to the baseband frequency. **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are√Ǭ†6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 √Ǭ†MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. 4f4926f603a6b760bcba99a1e8eb7640b78b9752 1057 1056 2012-04-12T20:59:05Z Akepley 16 /* About This Document */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]]<br /> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** If you're using doppler setting, then select "rest" next to the baseband frequency. **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32  MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Akepley]] 13:59, 12 April 2012 (PDT) 7abe3a92f92e89ce086480f623eee2716b60af1c 1058 1057 2012-04-12T21:01:08Z Akepley 16 /* About This Document */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]]<br /> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** If you're using doppler setting, then select "rest" next to the baseband frequency. **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32  MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 78333666acbfdab45bf35d7651c54c80c3e2ef4a 6 87 1035 2012-04-11T23:07:59Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 88 1036 2012-04-11T23:09:41Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 89 1041 2012-04-11T23:31:08Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 90 1059 2012-04-13T16:27:54Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 91 1060 2012-04-13T16:28:19Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 92 1061 2012-04-13T16:28:38Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 93 1062 2012-04-13T16:29:00Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 94 1063 2012-04-13T16:29:31Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 95 1064 2012-04-13T16:30:07Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 96 1065 2012-04-13T16:30:33Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 97 1066 2012-04-13T16:31:00Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 98 1067 2012-04-13T16:33:24Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 99 1068 2012-04-13T16:33:42Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 100 1069 2012-04-13T16:34:06Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 101 1070 2012-04-13T16:34:25Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 102 1071 2012-04-13T16:34:47Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 103 1072 2012-04-13T16:35:10Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 1073 1072 2012-04-13T16:35:35Z Akepley 16 uploaded a new version of &quot;[[File:New scan.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 104 1074 2012-04-13T16:35:57Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 105 1075 2012-04-13T16:37:42Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 106 1076 2012-04-13T16:38:03Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 107 1077 2012-04-13T16:38:19Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 108 1078 2012-04-13T16:38:37Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 109 1079 2012-04-13T16:39:04Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 110 1080 2012-04-13T16:39:56Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 111 1081 2012-04-13T16:40:31Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 1082 1081 2012-04-13T16:43:37Z Akepley 16 uploaded a new version of &quot;[[File:Scan loop.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 112 1083 2012-04-13T16:44:08Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 113 1084 2012-04-13T16:44:32Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 114 1085 2012-04-13T16:44:55Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 115 1086 2012-04-13T16:45:21Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 116 1087 2012-04-13T16:45:45Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 117 1088 2012-04-13T16:46:09Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 118 1089 2012-04-13T16:46:50Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 14 86 1090 1058 2012-04-13T16:47:59Z Akepley 16 /* OPT access and layout */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]]<br /> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** If you're using doppler setting, then select "rest" next to the baseband frequency. **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32  MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) d3ec8fa3fd28dfb15a407eb02a5c296d7e187f93 1091 1090 2012-04-13T16:48:26Z Akepley 16 /* Creating a Source List */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]]<br /> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** If you're using doppler setting, then select "rest" next to the baseband frequency. **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32  MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 4a369f3f5a753065783877d287bfafc5590bff5b 1092 1091 2012-04-13T16:50:32Z Akepley 16 /* Selecting a Complex Gain Calibrator */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]] You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** If you're using doppler setting, then select "rest" next to the baseband frequency. **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32  MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 8499d089eef9de6bd5b93474844fd8352d44038a 1093 1092 2012-04-13T16:51:12Z Akepley 16 /* Selecting a Complex Gain Calibrator */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]] You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). ** [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. ** [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** If you're using doppler setting, then select "rest" next to the baseband frequency. **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] **** [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32  MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]][[Image:dopset_output.png]][[Image:baseband_frequency_line_dopset.png]][[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 85673f12296a159f0fc3121dd0a7bd8591bc3473 1095 1093 2012-04-13T16:54:05Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]] You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** If you're using doppler setting, then select "rest" next to the baseband frequency. **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. **** [[Image:doppler_setting_source_selection.png]] *** If entering the line frequency manually, **** choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32  MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz. [[Image:dopset_input.png]] [[Image:dopset_output.png]] [[Image:baseband_frequency_line_dopset.png]] [[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 53201e9a6b40da1f0fcaa2d4e4ea71c8f2a2bdfd 1097 1095 2012-04-13T16:57:35Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]] You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** If you're using doppler setting, then select "rest" next to the baseband frequency. **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands. [[Image:baseband_frequency_line.png]] [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> *** If entering the line frequency manually, **** choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32  MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz. [[Image:dopset_input.png]] [[Image:dopset_output.png]] [[Image:baseband_frequency_line_dopset.png]] [[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 9b3feafa55d80ae9887c5145174ff1bb2a0954a6 1100 1097 2012-04-13T19:53:05Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]] You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range"[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** If you're using doppler setting, then select "rest" next to the baseband frequency. **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands. [[Image:baseband_frequency_line.png]] <br /> [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> *** If entering the line frequency manually, **** choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]] <br /> [[Image:dopset_output.png]] <br /> [[Image:baseband_frequency_line_dopset.png]] <br /> [[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) a03c2565c35b0463d3f30465f0e74755bcab7467 1101 1100 2012-04-13T19:56:07Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]] You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** If you're using doppler setting, then select "rest" next to the baseband frequency. **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.[[Image:baseband_frequency_line.png]] <br /> [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> *** If entering the line frequency manually, **** choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]] <br /> [[Image:dopset_output.png]] <br /> [[Image:baseband_frequency_line_dopset.png]] <br /> [[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 9237e77a287063d0d06b70e0d7117fe815c7598d 1102 1101 2012-04-13T19:59:25Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]] You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** Select "rest" next to the baseband frequency. **** You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.<br /> [[Image:baseband_frequency_line.png]] <br /> [[Image:subband_configuration_line.png]]<br /> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> *** If entering the line frequency manually, **** Choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]] <br /> [[Image:dopset_output.png]] <br /> [[Image:baseband_frequency_line_dopset.png]] <br /> [[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) f7325b5343433c3b0cb6a4db767ca602e07c970d 1103 1102 2012-04-13T20:02:16Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * Type the calibrator name in the search box at the top of the item manipulation region. [[Image:phase_calibrator_search.png]] * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.<br /> [[Image:phase_calibrator_select_group.png]] You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** Select "rest" next to the baseband frequency. **** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> *** If entering the line frequency manually, **** Choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]] <br /> [[Image:dopset_output.png]] <br /> [[Image:baseband_frequency_line_dopset.png]] <br /> [[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 313cc3d00d37b05a344ea58b0345379b6bbfde54 1104 1103 2012-04-13T20:03:41Z Akepley 16 /* Selecting a Complex Gain Calibrator */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever). [[Image:correlator_name_band.png]]<br /> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct. [[Image:subband_configuration.png]]<br /> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".[[Image:baseband_frequency.png]] *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** Select "rest" next to the baseband frequency. **** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> *** If entering the line frequency manually, **** Choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.[[Image:dopset_input.png]] <br /> [[Image:dopset_output.png]] <br /> [[Image:baseband_frequency_line_dopset.png]] <br /> [[Image:subband_configuration_line_dopset.png]] * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * Click on the "Pointing setups" group in the left hand side of the screen. [[Image:pointing_setups.png]]<br /> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) db77ed93593a4fc1d1121e34e7dfc88599cf1863 1105 1104 2012-04-13T20:06:02Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** Select "rest" next to the baseband frequency. **** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> *** If entering the line frequency manually, **** Choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]).The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following [[Image:opt_block_structure.png]]<br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * Click on the edit button for one of the source(s) that you will be observing in your scheduling block. [[Image:source_edit_button.png]] * In the editing window, click on the image tab to see the the elevation and azimuth curves for your source. * [[Image:source_image_tab.png]] [[Image:source_elevation_plot.png]]<br /> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * Edit the individual scans as needed. [[Image:bulk_scan_creation.png]]<br /> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. [[Image:scan_loop.png]]<br /> If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: [[Image:bulk_scan_edit_select.png]]<br /> To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source. [[Image:bulk_scan_edit_change.png]]<br /> * Now click "Update" * A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters. [[Image:bulk_scan_edit_confirm.png]]<br /> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 20b86520afec219bb76d46eaed4098780df17b1f 1106 1105 2012-04-13T20:09:08Z Akepley 16 /* Putting It All Together: Creating a Scheduling Block */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** Select "rest" next to the baseband frequency. **** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> *** If entering the line frequency manually, **** Choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: [[Image:opt_block_structure.png]] The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 57961282939aa9aa50b4531b697f511fae780462 1107 1106 2012-04-13T20:23:21Z Akepley 16 /* Putting It All Together: Creating a Scheduling Block */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** Select "rest" next to the baseband frequency. **** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> *** If entering the line frequency manually, **** Choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: [[Image:opt_block_structure.png]] The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block. [[Image:lst_start_range.png]] You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 2c047339b8e6fc45688fc63a8d552d210e203cbe 1108 1107 2012-04-13T20:24:44Z Akepley 16 /* Putting It All Together: Creating a Scheduling Block */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * The Observation Preparation section where you prepare your scheduling blocks, * The Sources section where you put together your source lists, and * The Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** Select "rest" next to the baseband frequency. **** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> *** If entering the line frequency manually, **** Choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 4c76c05412b6fa556dbe65fcd7bf12ed37f17d4e 1109 1108 2012-04-13T21:53:12Z Akepley 16 /* OPT access and layout */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> *** There are also default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. ** For spectral line observations: *** If doppler setting, **** Select "rest" next to the baseband frequency. **** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> *** If entering the line frequency manually, **** Choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 86faa4b901087f11f5601a3748faa8c674368c06 1110 1109 2012-04-13T21:56:49Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> *** If doppler setting, **** Select "rest" next to the baseband frequency. **** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> **** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> *** If entering the line frequency manually, **** Choose "sky" in the Frequencies section. **** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool **** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. **** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. **** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) b0f95cac21ffb72f7ef37521ea5a620771c37691 1111 1110 2012-04-13T22:00:07Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the line frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 15a2391589aba162aef169c7e5c44d539c1c7905 1112 1111 2012-04-13T22:01:02Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. It's a good idea to use this number to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 0f54cb824b9cf52ac084928ab65887e50cf2ffe5 1113 1112 2012-04-13T22:03:06Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * <p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST restrictions are a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 89ce3be70972418868ceb0ab01a3ad482363174d 1114 1113 2012-04-13T22:03:42Z Akepley 16 /* Putting It All Together: Creating a Scheduling Block */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to akepley@nrao.edu. A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * <p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 1d39652e00407ec63aa2d5d3fefd63f6f3ce76dc 1115 1114 2012-04-13T23:38:01Z Akepley 16 /* Introduction */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. [[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * <p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) f16004dc421caeaa696a7f7e4a8ed3857ba657c2 1116 1115 2012-04-13T23:38:53Z Akepley 16 /* Selecting a Bandpass Calibrator */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). A scheduling block defines a complete set of EVLA observations including your source observations and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * <p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 8b3945a3732b0094cc63a168a20ba16584691a00 1117 1116 2012-04-16T20:57:43Z Akepley 16 /* Introduction */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * <p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) f4d6c0edfdd5d48f9c68b9c54845981c1f7b92c6 1118 1117 2012-04-16T20:59:23Z Akepley 16 /* Introduction */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * <p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 591a3c72dfae0d8bb7fe2519cb701f93a39ea0aa 1119 1118 2012-04-16T21:00:44Z Akepley 16 /* Overview of the Scheduling Block Creation Process */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * <p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) f21f34e9deb27caad36a90b6857fa7ba20c37a60 1120 1119 2012-04-16T21:01:30Z Akepley 16 /* Overview of the Scheduling Block Creation Process */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * <p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 591a3c72dfae0d8bb7fe2519cb701f93a39ea0aa 1121 1120 2012-04-16T21:03:43Z Akepley 16 /* OPT access and layout */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source position and name. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. Your source names need to be less than 14 characters for your data to be compatible with AIPS. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (generally &lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, you can get information on that source taken from VLA calibrator manual. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * <p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) bff35ce5766276cf95befda97f1653ef24405e88 1122 1121 2012-04-16T21:05:52Z Akepley 16 /* Creating a Source List */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * <p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 17fdc7e8f0e15a4811ac55d71a0011e350084c84 1123 1122 2012-04-16T21:08:03Z Akepley 16 /* Selecting a Flux Calibrator */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/structure column, you can get the VLA calibrator manual information for that calibrator. Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * <p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) d06af3559a576f2d9544b8f6087e35a2b3ae46d5 1124 1123 2012-04-16T21:10:04Z Akepley 16 /* Selecting a Flux Calibrator */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * <p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 11886180c606a5625f67beba051dfcb73ee93c32 6 119 1094 2012-04-13T16:52:03Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 120 1096 2012-04-13T16:54:38Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 121 1098 2012-04-13T19:39:17Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 122 1099 2012-04-13T19:39:59Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 123 1125 2012-04-16T21:10:56Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 14 86 1126 1124 2012-04-16T21:12:03Z Akepley 16 /* Selecting a Flux Calibrator */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add these configuration to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. After creating a new OSRO configuration, the OSRO configuration editor will appear in the main editing window. We will work our way through the instrument configuration settings starting at the top of the instrument configuration editor. * In the first section of the instrument configuration page, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 9GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> * <p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) a32e1992edf69fbc94e0e28e907bf65c1344bac3 1129 1126 2012-04-16T21:30:20Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. The observing conditions may change significantly over the course of longer blocks (&gt;4 hrs) while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block into this consideration. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the appropriate values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and phase calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 0387580be6caf3c5e03d79a213cd168172477c37 1130 1129 2012-04-16T21:38:13Z Akepley 16 /* Putting It All Together: Creating a Scheduling Block */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after its been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) f07498e94c85953da09dc12ba5e8640eb16ae7bf 1188 1130 2012-04-26T16:35:48Z Akepley 16 /* Submitting Your Scheduling Block */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * Create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 1e56f6b7b18fb57f20c45f3afbe0c790a296fa51 1215 1188 2012-05-23T20:28:54Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 30min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 005ce2abfc1eab281ee072d9c7ec43a772b6f066 6 112 1127 1083 2012-04-16T21:14:19Z Akepley 16 uploaded a new version of &quot;[[File:Skymap option.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 91 1128 1060 2012-04-16T21:24:06Z Akepley 16 uploaded a new version of &quot;[[File:Baseband frequency.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 0 48 1186 1031 2012-04-26T03:01:39Z Jott 8 wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. The JVLA sesntitivy of spectral lines, however, is almost the same as with the VLA. Where the JVLA capabilities really make the difference is in the number of channels and flexibility of the spectral line setups. The JVLA can * observe multiple lines at the same time * deliver continuous spectral coverage up to a full width 8GHz * access 1GHz or 2GHz chinks of each receiver band and place correlator subbands on it * use up to 64 subbands at a time which almost act like 64 independent correlators, independently tunable, with different spectral bandwidths, different channel numbers, and different polarization product coverage * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage] and the [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips & tricks on using CASA and the visualization tools, such as the [http://casa.nrao.edu/CasaViewerDemo/casaViewerDemo.html CASA viewer] that are designed to display spectral data cubes == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 441a8e9e1236f1dcd5e8e9f374c5c3db6e93746a 1230 1186 2012-06-07T22:09:25Z Jott 8 /* JVLA Spectral Line Observing */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage] and the [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips & tricks on using CASA and the visualization tools, such as the [http://casa.nrao.edu/CasaViewerDemo/casaViewerDemo.html CASA viewer] that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu-\nu_0}{\nu_0}\,\,c = \frac{\lambda_0-\lambda}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 4e7789220e6baa9a68227b9aaebccc1d7a976ee7 1231 1230 2012-06-07T23:20:45Z Jott 8 /* Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions to high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage] and the [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips & tricks on using CASA and the visualization tools, such as the [http://casa.nrao.edu/CasaViewerDemo/casaViewerDemo.html CASA viewer] that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 3cc398cb3c1203b1c140a78fd227f9ba76f7a9cd 1232 1231 2012-06-07T23:37:35Z Jott 8 /* JVLA Spectral Line Observing */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy] and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 3b8c9bf889d60d3d66f65b4995b3b931d31d5d55 1233 1232 2012-06-07T23:40:09Z Jott 8 /* Line Rest Frequencies */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 82bce61fbd06f15eea0764b8efbd19a545b968c0 1234 1233 2012-06-07T23:41:45Z Jott 8 /* Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity is is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 65c5b34f02b2f72d0f5fb88922846018c0534ddc 1235 1234 2012-06-07T23:46:19Z Jott 8 /* Velocity Frames */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This is called the ''relativistic velocity''. This equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is possible to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> The redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 031b94a97266dfae37fcd4b0e24a3779c573345e 1236 1235 2012-06-07T23:49:10Z Jott 8 /* Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame to which the velocities are measured to. There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceding ones: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen to what is appropriate for the science. Three frames are commonly used: * '''Topocentric''' this is the frame that the sky (observing frequency) uses. It is also the standard for visibilities in the measurement set * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. The standard used in almost all cases is LSRK and most likely the older LSR naming is identical to the mode modern LSRK definition. * '''Barycentric''' is a common frame, too and has virtually replaced the older heliocentric standard. Given the small difference between them, they were frequently used interchangeably. A full list of reference frames that CASA supports is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric. == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. f6cf5435e458f3625e228b8876c83f9ace4da95c 1237 1236 2012-06-08T00:01:35Z Jott 8 /* Velocity Frames */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 3f2197e77b9a635075e5328bd2b80fe03ee5ddfd 1239 1237 2012-06-20T06:16:09Z Jott 8 wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. d68d2c6af6f26c6e038480eb548b712f05c63fa7 1240 1239 2012-06-20T06:18:03Z Jott 8 /* Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c = v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 129b98cbc90d8fa303a185d9904b2d920866554e 1241 1240 2012-06-20T06:18:44Z Jott 8 /* Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope operates typical at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during a typical observation campaign. Within a single observation, the rotation of the earth dominates and the line may shift up to ~0.5 km/s (see above). observing campaigns that span longer times, may see spectral lines to shift in frequency by up to 30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds to <math>x</math> km/s for the line at a wavelegth of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1 MHz corresponds to about 7 km/s in velocity. </i> Using this rule of thumb, a line may shift about up to 5MHz in Q-band and up to 0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a different, but fixed sky frequency for each observation. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (0.5 km/s max). This small shift is corrected again by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. A non-variable sky frequency delivers also a more robust calibration and system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. This will require that spectral features need to be sampled with at least 4 channels to be correctly reproduced. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 95a8a86797f1d66d992d21eace11c73b410df172 1243 1241 2012-06-20T06:34:08Z Jott 8 /* Doppler Correction */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm 0.5</math>km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm 30</math>km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm5</math>MHz in Q-band and by up to <math>0.15</math>MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm 30</math>km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <mtah>\pm0.5</math>km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 0ff1695d13ab876f12b740651250619f3777d015 1244 1243 2012-06-20T06:36:13Z Jott 8 /* Doppler Correction */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design could show the Gibbs phenomenon relatively prominently and frequently Hanning smoohting was applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or rfi sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there's no need for data size reduction via Hanning smoothing anymore. As a consequence, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. a24e622a1afc98cadfc1ebb32ef5928ade05f349 1245 1244 2012-06-20T06:40:58Z Jott 8 /* Gibbs Phenomenon and Hanning smoothing */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity Calculation == The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA exposure calculator]. This JAVA tool allows you to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the velocity of the resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 53c720a119c13a4a43d16a180a4b9b013c71db66 1247 1245 2012-06-20T06:47:20Z Jott 8 /* Sensitivity Calculation */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator The JVLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. a5da7c63985c723d6ece3435281cc308646b884e 1248 1247 2012-06-20T06:48:31Z Jott 8 /* Sensitivity/Exposure Time Calculation */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|600px|thumb|right|The JVLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. c9c356f83251c65096325fdb5e9019a287f6ddf3 1250 1248 2012-06-20T06:51:44Z Jott 8 /* Sensitivity/Exposure Time Calculation */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The JVLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == The WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and during that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations and are the most fundamental spectral ranges for any observations with WIDAR. In the 8-bit mode, WIDAR features two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. As a second step, the basebands enter digital filters, 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|600px|thumb|right|WIDAR correlator baseband with subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After the analog filter that define the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We have a couple of tools that help, please check the spectral line section of the [https://safe.nrao.edu/wiki/bin/view/EVLA/RSROObservingPreparationGuidelines JVLA RSRO Observing Preparation Guidelines] * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To get more channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|400px|thumb|right|Baseband with 128MHz suckouts]] rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges. [[File:BlankFieldRMS.interlace.png|400px|thumb|right|Shifted baseband setup to substitute suckout channels]] rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted. === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. f8f8dff2fd54e62bddd9c1dbecc47057149c0981 1251 1250 2012-06-20T07:05:32Z Jott 8 /* The WIDAR Correlator */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The JVLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|200px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|200px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands. Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 216 channelf for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 2896370d3c5c2ac744a722f23bbacf1046aacd3d 1252 1251 2012-06-20T07:07:53Z Jott 8 /* Narrow Subbands */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see xxxxx): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The JVLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|200px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|200px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 2e46e68815a62104ab734932bef4d4c033e8ae94 1253 1252 2012-06-20T07:09:58Z Jott 8 /* JVLA Spectral Line Observing */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The JVLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|200px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|200px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == === The Proposal Submission Tool (PST) === === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 1393fa4f7a9c6bec6d0cfd8d94ea36368049f361 1256 1253 2012-06-20T07:23:29Z Jott 8 /* Planning and Setup */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The JVLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|200px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|200px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST, accessible via [http://my.nrao.edu my.nrao.edu]) includes a spectral setup tool. [[File:PST-wide.png]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 4106e046ecd5b76b6475c197378fd35748c0373e 1257 1256 2012-06-20T07:24:34Z Jott 8 /* Fixed 128MHz Subbands and 128 MHz "Suckouts" */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The JVLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST, accessible via [http://my.nrao.edu my.nrao.edu]) includes a spectral setup tool. [[File:PST-wide.png]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 774a59f402a9894d2be9860d35bbe8958fbede4e 1258 1257 2012-06-20T07:27:11Z Jott 8 /* The Proposal Submission Tool (PST) */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The JVLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, only one baseband can be below 32GHz and that must be BD * The A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz separation max. For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST, accessible via [http://my.nrao.edu my.nrao.edu]) includes a spectral setup tool. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. b14b89ddc05d1ed740177db8371ac6f58757eb01 1259 1258 2012-06-20T07:28:51Z Jott 8 /* WIDAR Tuning Restrictions */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The JVLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. in OPT, this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to go to wider bandwidths in each subband but yet maintain a high number of channels. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each subbands, quadrupling the number of channels from 128 to 512, or reducing the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth essentially providng a very high spectral resolution for smaller subbands. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST, accessible via [http://my.nrao.edu my.nrao.edu]) includes a spectral setup tool. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. a16b0dc3023e0365f6de38da20f8642696c53dd3 1260 1259 2012-06-20T07:33:24Z Jott 8 /* Baselineboard Stacking */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The JVLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to get more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is not available right now in OPT and JVLA staff needs to be contacted. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST, accessible via [http://my.nrao.edu my.nrao.edu]) includes a spectral setup tool. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 285228d3d8eef183c0d082a8e561c4d35796370d 1261 1260 2012-06-20T07:34:29Z Jott 8 /* Recirculation */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The JVLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 75MB/s. The OPT instrument configuration calculates them based on the spectral line setup and the limit of 65 MB/s should not be exceeded. == Planning and Setup == For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST, accessible via [http://my.nrao.edu my.nrao.edu]) includes a spectral setup tool. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 702994dd95a6e3d889bb4bb9cde6f8b3af8fcaae 1262 1261 2012-06-20T07:35:54Z Jott 8 /* Data Rate Limits */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The JVLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup == For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST, accessible via [http://my.nrao.edu my.nrao.edu]) includes a spectral setup tool. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during an observation is run. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. You may need to register at my.nrao.edu if you do not yet have an account. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequecnsy setup as shown in the figures. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behaviour of the receiver sensitivity at the edges. If the frequency to be observed is close to the edges of the receiver, one may check if the next higher or lower frequency receiver is more suitable * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', one can choose 2 basebands, each with a width of 1GHz. The center frequencies will go into the AOCO box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', one can choose 4 basebands, each 2 GHz wide. Here the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional freqeuny restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation. On the other hand, time smearing effects, rfi excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, in case that a couple of them are being worked on. Also remember that the maximum supported data rate is currently 75MB/s. If you require higher data rates, please contact NRAO staff. * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each of the basebands selected. Under each tab one can now select the individual subbands. Up to 64 subbands are available: click ''Add subband'' to create a subband setting, where one can select the frequency range from a "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. Currently, the subbands are not independently tuneable yet (it is under commissioning) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT (??? rest frequencies for OSRO ???). The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: please contact NRAO staff if you need this setup * Independently tuning the subbands: Please contact NRAO staff for this feature [[File:OPT-config1.png|400px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|400px|thumb|right|OPT - Instrument Configuration: Subband Settings]] ==== Doppler Setting ==== Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation and the frequency will stay the same for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode, RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by max. 30 km/s over a year, 0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use, first select ''Rest'' in the baseband frequency setup section of the OPT/ICT. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. In the future this feature may be available for each subband separately but for now it will shift the baseband frequency and shift all subbands along with it. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * Of the feature is narrow one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, a option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. FHowever, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined on by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used is the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be advisable to subtract the continuum using line-free channels in the image cube. Both methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model of the continuum. If a the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be obtained and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy srong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. ==== Zeeman Observing ==== please contact NRAO local staff if you plan to perform Zeeman observations. This mode is currently in the commissioning phase. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. a2700b18b19cd213dc8219698f6b66ae8809ff0d 1265 1262 2012-06-20T08:25:26Z Jott 8 /* Planning and Setup */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The JVLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of JVLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of JVLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. 072f15da10fe92e9415324e17f158d6a9ab55384 1266 1265 2012-06-20T08:26:11Z Jott 8 /* Spectral Line Observing */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The JVLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of JVLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of JVLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. --> c3bf1dbda6ce60e8b557e85f72a67e7c5fe4f7db 1267 1266 2012-06-20T08:35:50Z Jott 8 /* Doppler Setting */ wikitext text/x-wiki The new JVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == JVLA Spectral Line Observing == This guide is to help understand and setting up spectral line observations at the JVLA. The new, wide bandwidths of the JVLA allow users to observe up to 8GHz spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the JVLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final JVLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configures in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for JVLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on JVLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for JVLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi JVLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The JVLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the JVLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the JVLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The JVLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure JVLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions JVLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the JVLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the JVLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The JVLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of JVLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss JVLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the JVLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of JVLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the JVLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the JVLA (last paragraph) The JVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for JVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the JVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the JVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the JVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the JVLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive JVLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process JVLA spectra line data. --> 2d1672f1dd7dd660ae3c1076bd179611de7497cd 1268 1267 2012-06-27T05:07:40Z Jott 8 wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 87f9b9e428954d0869f9f557207bec910f2ce894 0 125 1222 2012-05-24T15:46:51Z Jott 8 Created page with "'''REPLACED ON 23 MAY 2012 BY https://science.nrao.edu/facilities/evla/docs/manuals/oss''' '''WE RECOMMEND YOU UPDATE YOUR BOOKMARKS ACCORDINGLY''' ---- [[Observational Summary..." wikitext text/x-wiki '''REPLACED ON 23 MAY 2012 BY https://science.nrao.edu/facilities/evla/docs/manuals/oss''' '''WE RECOMMEND YOU UPDATE YOUR BOOKMARKS ACCORDINGLY''' ---- [[Observational Summary 29-Nov-2011 frozen as of 24-May-2012]] '''The EVLA Observational Status Summary''' ''Version date: November 29, 2011'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget at the end of 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || align='center'| 10 μJy || align='center'| 1 μJy || align='center'| 10 |- | Maximum BW in each polarization || align='center'| 0.1 GHz || align='center'| 8 GHz || align='center'| 80 |- | Number of frequency channels at max. BW || align='center'| 16 || align='center'| 16,384 || align='center'| 1024 |- | Maximum number of freq. channels || align='center'| 512 || align='center'| 4,194,304 || align='center'| 8192 |- | Coarsest frequency resolution || align='center'| 50 MHz || align='center'| 2 MHz || align='center'| 25 |- | Finest frequency resolution || align='center'| 381 Hz || align='center'| 0.12 Hz || align='center'| 3180 |- | Number of full-polarization sub-correlators || align='center'| 2 || align='center'| 64 || align='center'| 32 |- | Log (Frequency Coverage over 1–50 GHz) || align='center'| 22% || align='center'| 100% || align='center'| 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date !Status or Actual Date |- | Installation of EVLA correlator subset for early science || align='center'| 2010 Q1 || align='center' | 2012 Q1 |- | Shared Risk Observing begins || align='center'| 2010 Q1 || align='center' | 2012 Q2 |- | Last antenna retrofitted || align='center'| 2010 Q2 || align='center' | 2010 Q2 |- | Full EVLA correlator installation || align='center'| 2011 Q2 || align='center' | 2011 Q2 |- | Last receiver installed || align='center'| 2012 Q4 || align='center' | on schedule |- |} == VLA to EVLA Transition == The year 2010 was extremely exciting for the EVLA. The correlator that was the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from '''A'''→'''B'''→'''C'''→'''D'''→'''A''' to '''D'''→'''C'''→'''B'''→'''A'''→'''D''', in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. The last VLA antenna was retrofitted to EVLA specifications in May 2010. During 2011 the WIDAR correlator was put into full observing mode with commissioning and Resident Shared Risk Observing starting in early 2011. By the end of 2011, up to 2 GHz of bandwidth was provided to the the general public along with 2 tunable bands, each with 8 spectral windows (64 to 256 channels) at S, Ku and X bands. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it has not yet been commissioned. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted '''D''', '''C''', '''B''', and '''A''' respectively. In addition, there are 3 “hybrid” configurations labelled '''DnC''', '''CnB''', and '''BnA''', in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle was modified during 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule for 2011, 2012 and 2013 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Approximate EVLA Configuration Schedule for 2011-2012''' ! Year ! Jan-Apr ! May-Aug ! Sep-Dec |- | align='center'| 2011 || align='center'| '''B''' || align='center'| '''A''' || align='center'| '''D''' |- | align='center'| 2012 || align='center'| '''C''' || align='center'| '''B''' || align='center'| '''A''' |- | align='center'| 2013 || align='center'| '''D''' || align='center'| '''C''' || align='center'| '''B''' |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 2012 (a full '''D'''→'''A''' configuration cycle) are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science is provided by two programs for outside users and one for EVLA commissioning staff. All early science programs are peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period thus involves an element of risk associated with the large stepwise increases in throughput bandwidth that are offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program provides early science capabilities to the general user community. These capabilities initially provided a maximum 256 MHz bandwidth that increased to 2 GHz for the '''D''' configuration in mid-2011 and will increase further to 8 GHz at the end of 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All EVLA antennas are outfitted with either EVLA or “interim” L, EVLA C, VLA X, EVLA K, EVLA Ka, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of July 2012, 25 of the EVLA antennas will be outfitted with S-band receivers, 23 EVLA antennas will have Ku-band receivers, and 22 will have new EVLA-style X-band receivers. Figure 1 shows the expected installation rate of final EVLA receiver systems for the rest of the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers, expected to be commissioned by the end of 2012. Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:RcvrAvailDec12.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of final EVLA receivers during the calendar year 2012 until the end of the EVLA Construction Project on December 31, 2012. Interim receivers with reduced frequency coverage or polarization purity are available at some bands in addition to those plotted (see Table 4).]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that are available in September 2011 and January 2012, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, Jan 2012 ! colspan="2" | Receiver availability, Jul 2012 |- ! ! GHz ! Final EVLA ! EVLA+VLA/interim ! Final EVLA ! EVLA+VLA/interim |- | 400 cm (4-band) || align='center'| 0.062–0.078 || align='center'| - || align='center'| - || align='center'|- || align='center'|- |- | 20 cm (L) || align='center'| 1.0–2.0 || align='center'| 19 || align='center'| 27 || align='center'|24 || align='center'|27 |- | 13 cm (S) || align='center'| 2.0–4.0 || align='center'| 21 || align='center'| 21 || align='center'|25 || align='center'|25 |- | 6 cm (C) || align='center'| 4.0–8.0 || align='center'| 27 || align='center'| 27 || align='center'|27 || align='center'|27 |- | 3 cm (X) || align='center'| 8.0–12.0 || align='center'| 16 || align='center'| 27 || align='center'|22 || align='center'|27 |- | 2 cm (Ku) || align='center'| 12.0–18.0 || align='center'| 18 || align='center'| 18 || align='center'|23 || align='center'|23 |- | 1.3 cm (K) || align='center'| 18.0–26.5 || align='center'| 27 || align='center'| 27 || align='center'|27 || align='center'|27 |- | 1 cm (Ka) || align='center'| 26.5–40.0 || align='center'| 27 || align='center'| 27 || align='center'|27 || align='center'|27 |- | 0.7 cm (Q) || align='center'| 40.0–50.0 || align='center'| 27 || align='center'| 27 || align='center'|27 || align='center'|27 |- |} :Note: The "EVLA+VLA/interim" columns give the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers (see Figure 1). == Open Shared Risk Observing (OSRO) == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program has extended this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program obtain good quality data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing (RSRO) == The WIDAR correlator and the EVLA provides a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program is now expected to run through the end of 2012, with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. Full operations of the EVLA will begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program for the upcoming '''D'''→'''A''' configuration cycle, Sep 2011 through Dec 2012. Users interested in participating in the RSRO program should refer to the web page at https://science.nrao.edu/facilities/evla/early-science/rsro for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24 || 80 || 260 || 850 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka) || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800 || 2200 || 20000 || 20000 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka) || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations ('''DnC''', '''CnB''', '''BnA''') should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., '''DnB''', or '''CnA''') is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. For the '''C''' configuration an antenna from the middle of the north arm is moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Note that although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration. ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are insufficient EVLA 8–12 GHz receivers available yet to determine system performance across the band. A project with the goal of doubling the longest baseline available in the '''A''' configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the ''SEFD'' at some fiducial EVLA frequencies. [[File:SEFD.png|none|frame|Figure 2: ''SEFD'' for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the Ku, K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: ''SEFD''s and '''D'''-Configuration Confusion Limits''' !Frequency !''SEFD'' !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || align='center'| 420 || align='center'| 89 |- | 3.0 GHz (S) || align='center'| 370 || align='center'| 14 |- | 6.0 GHz (C) || align='center'| 310 || align='center'| 2.3 |- | 10.0 GHz (X) || align='center'| 250 || align='center'| negligible |- | 15 GHz (Ku) || align='center'| 350 || align='center'| negligible |- | 22 GHz (K) || align='center'| 560 || align='center'| negligible |- | 33 GHz (Ka) || align='center'| 730 || align='center'| negligible |- | 45 GHz (Q) || align='center'| 1400 || align='center'| negligible |} :Note: ''SEFD''s at Ku, K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in '''D''' configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temperature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted above assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. '''At source elevations greater than 80 degrees (zenith angle < 10 degrees), source tracking becomes difficult; it is recommended to avoid such source elevations during the observation preparation setup'''. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in Right Ascension (RA). The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where ''T''_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and ''S'' in mJy per beam, the constant ''F'' depends only upon array configuration and has the approximate value ''F'' = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for ''S''. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of ''F'' calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where ''D'' is the distance to the galaxy in Mpc, and ''S∆V'' is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different “baseband” frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” Currently, a maximum of 1.024 GHz can be correlated for each IF pair (see [[#Correlator Configurations|Correlator Configurations]]), for a total maximum bandwidth of approximately 2 GHz. The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. At X-band a number of the antennas will continue to have the old narrow-band VLA X-band receivers until their retrofit is complete at the end of the EVLA construction project. As of December 2010 there is not a sufficient number of EVLA-style X-band receivers in the array to evaluate either the system performance or the radio frequency interference environment throughout the EVLA X-band tuning range of 8-12 GHz. A total bandwidth of 800 MHz equivalent to that of the VLA receivers (8.0-8.8 GHz) should be assumed for the purposes of sensitivity calculations at X-band for the present. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || align='center'| 1.0–2.0 || align='center'| 1.25¹ || align='center'| 1.75¹ |- | 13 cm (S) || align='center'| 2.0–4.0 || align='center'| 2.5 || align='center'| 3.5 |- | 6 cm (C) || align='center'| 4.0–8.0 || align='center'| 5.0 || align='center'| 6.0 |- | 3 cm (X) || align='center'| 8.0–12.0 || align='center'| 8.5² || align='center'| 9.5² |- | 2 cm (Ku) || align='center'| 12.0–18.0 || align='center'| 13.5 || align='center'| 14.5 |- | 1.3 cm (K) || align='center'| 18.0–26.5 || align='center'| 20.7 || align='center'| 21.7 |- | 1 cm (Ka) || align='center'| 26.5–40.0 || align='center'| 31.5 || align='center'| 32.5 |- | 0.7 cm (Q) || align='center'| 40.0–50.0 || align='center'| 40.5 || align='center'| 41.5 |} :Notes: :: 1. This default frequency set-up for L-band comprises two 512 MHz basebands (each with 8 subbands of 64 MHz) to cover the entire 1-2 GHz of the L-band receiver. :: 2. Many of the antennas continue to have the old narrow-band VLA receivers (see Figure 1), for which a total bandwidth of 800 MHz should be assumed (8.0-8.8 GHz). The RFI environment of the default tunings has not yet been evaluated. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 in [[#Documentation|Documentation]] for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! (∆ν/ν₀)*(θ₀/θ_HPBW) ! Peak ! Width |- | 0.0 || align='center'| 1.00 || align='center'| 1.00 |- | 0.50 || align='center'| 0.95 || align='center'| 1.05 |- | 0.75 || align='center'| 0.90 || align='center'| 1.11 |- | 1.0 || align='center'| 0.80 || align='center'| 1.25 |- | 2.0 || align='center'| 0.50 || align='center'| 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 (in [[#Documentation|Documentation]]) considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || align='center'| 2.1 || align='center'| 4.8 || align='center'| 6.7 |- | '''B''' || align='center'| 6.8 || align='center'| 15.0 || align='center'| 21.0 |- | '''C''' || align='center'| 21.0 || align='center'| 48.0 || align='center'| 67.0 |- | '''D''' || align='center'| 68.0 || align='center'| 150.0 || align='center'| 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 (http://www.aoc.nrao.edu/evla/geninfo/memoseries/evlamemo67.pdf) for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second, for '''A''' configuration. For '''B''', '''C''', and '''D''' configurations the minimum integration time is 3 seconds, unless a shorter integration time is explicitly requested and justified. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 14.4 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 51.8 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 1.2 TB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool (https://science.nrao.edu/facilities/evla/data-archive/evla) will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on disk drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the '''D''' configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum at L-band. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots and tables of known RFI are available online, at http://science.nrao.edu/evla/observing/RFI/; plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the '''D''' configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in '''A''' configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in '''A''' configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List (http://www.vla.nrao.edu/astro/calib/manual/) should be used. The positions of these sources are taken from lists published by the United States Naval Observatory (USNO). == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1 in [[#Documentation|Documentation]]. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth smearing and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its '''D''' configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP (Wilkinson Microwave Anisotropy Probe) flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby calibrator with an "S" code in the calibrator database, and a more distant calibrator with a "P" code, the nearby calibrator is usually the better choice (see http://www.vla.nrao.edu/astro/calib/manual/key.html for a description of calibrator codes). :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173 (http://www.vla.nrao.edu/memos/sci/). These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API may be found at http://www.vla.nrao.edu/astro/guides/api/. Plots of current/historical data can be found at: https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi Characteristic seasonal averages are represented below: {| border="1" align="center" |+ '''Table: Seasonal API/wind values at the EVLA''' !Month !API (night) [deg] !API (median) [deg] !API (day) [deg] !Wind (night) [m/s] !Wind (median) [m/s] !Wind (day) [m/s] |- | [[Media:APIwind_January.png| January]] || 2.3 || 2.8 || 3.6 || 1.6 || 1.9 || 2.3 |- | [[Media:APIwind_February.png| February]] || 2.9 || 3.4 || 4.5 || 4.0 || 4.3 || 4.5 |- | [[Media:APIwind_March.png| March]] || 2.8 || 3.7 || 5.5 || 3.4 || 3.9 || 4.7 |- | [[Media:APIwind_April.png| April]] || 3.3 || 4.5 || 6.2 || 5.3 || 5.5 || 5.8 |- | [[Media:APIwind_May.png | May]] || 2.9 || 4.6 || 6.7 || 2.6 || 3.2 || 3.7 |- | [[Media:APIwind_June.png| June]] || 3.8 || 5.5 || 7.4 || 2.5 || 3.9 || 6.3 |- | [[Media:APIwind_July.png| July]] || 6.2 || 8.3 || 10.5 || 2.9 || 2.9 || 3.0 |- | [[Media:APIwind_August.png| August]] || 5.4 || 7.1 || 11.3 || 1.7 || 2.3 || 3.0 |- | [[Media:APIwind_September.png| September]] || 5.2 || 6.6 || 8.8 || 2.3 || 3.0 || 3.6 |- | [[Media:APIwind_October.png| October]] || 4.2 || 5.3 || 7.4 || 2.3 || 2.9 || 3.7 |- | [[Media:APIwind_November.png| November]] || 2.6 || 3.0 || 4.0 || 1.2 || 2.5 || 1.6 |- | [[Media:APIwind_December.png| December]] || 2.8 || 3.2 || 4.1 || 1.2 || 1.6 || 2.7 |- |} Click on the Month links above to see plots of phase and wind speed versus time. :Note: day indicates sunrise to sunset values; night indicates sunset to sunrise values. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/<math>{\rm{m}^2}</math>, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. For more details see: https://science.nrao.edu/facilities/evla/early-science/polarimetry == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. For the Open Shared Risk Observing (OSRO) program available to the community during the period Sep 2011 through Jan 2013 are offering two independently tunable basebands, where each baseband has up to eight sub-bands. Possible sub-band widths are 128 MHz, 64 MHz, 32 MHz, all the way down in factors of 2 to 0.03125 MHz. All sub-bands must have the same bandwidth and channelization in both basebands, and be contiguous in frequency within each baseband. We will offer three different OSRO modes: full polarization, dual polarization, and single polarization, with 64, 128, and 256 channels per sub-band, respectively. There is always the possibility during offline processing to smooth in frequency to reduce dataset sizes or to improve spectral response. Starting with the '''D'''-configuration in September 2011, we have been providing options for configuring WIDAR for OSRO in the following three ways: :1. “OSRO Full Polarization”: Four polarization products. This configuration offers 4 polarization products for each sub-band, each of which has 128 MHz bandwidth with 64 channels. It is possible to decrease the subband bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in the following table: {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO Full Polarization)''' ! Sub-band BW (MHz) ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 64 || 1000 || 300 || 19,200 |- | 32 || 64 || 500 || 150 || 9,600 |- | 16 || 64 || 250 || 75 || 4,800 |- | 8 || 64 || 125 || 37.5 || 2,400 |- | 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 64 || 31.25 || 9.4 || 600 |- | 1 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO Dual Polarization”: Two polarization products. This configuration offers 2 polarization products for each sub-band, each of which has 128 MHz bandwidth with 128 channels. It is possible to decrease the sub-band bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in the following table. {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for dual polarization (OSRO Dual Polarization)''' ! Sub-band BW (MHz) ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 128 || 1000 || 300/ν (GHz) || 38,400/ν (GHz) |- | 64 || 128 || 500 || 150 || 19,200 |- | 32 || 128 || 250 || 75 || 9,600 |- | 16 || 128 || 125 || 37.5 || 4,800 |- | 8 || 128 || 62.5 || 19 || 2,400 |- | 4 || 128 || 31.25 || 9.4 || 1,200 |- | 2 || 128 || 15.625 || 4.7 || 600 |- | 1 || 128 || 7.813 || 2.3 || 300 |- | 0.5 || 128 || 3.906 || 1.2 || 150 |- | 0.25 || 128 || 1.953 || 0.59 || 75 |- | 0.125 || 128 || 0.977 || 0.29 || 37.5 |- | 0.0625 || 128 || 0.488 || 0.15 || 18.75 |- | 0.03125 || 128 || 0.244 || 0.073 || 9.375 |} :3: "OSRO Single Polarization": One polarization product (new for OSRO observing). It offers 1 polarization product for each sub-band, each of which has 128 MHz bandwidth with 256 channels. It is possible to decrease the sub-band bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in the following table. {| border="1" align="center" |+ '''Table 14: Correlator capabilities per sub-band for single polarization (OSRO Single Polarization)''' ! Sub-band BW (MHz) ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 256 || 250 || 75 || 19,200 |- | 32 || 256 || 125 || 37.5 || 9,600 |- | 16 || 256 || 62.5 || 19 || 4,800 |- | 8 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities are being provided with integration times no shorter than 1 second in '''A''' configuration (3 seconds in '''B'''/'''C'''/'''D''' configurations), and Doppler setting will be available with these correlator configurations. If it is likely that the data will need to be resampled spectrally in order to Doppler track to a line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna (e.g., Y1) modes, have not yet been commissioned and are not yet available to the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing "Faint Images of the Radio Sky at Twenty-centimeters survey(FIRST, http://www.cv.nrao.edu/first/) or the Co-Ordinated Radio 'N' Infrared Survey for High-mass star formation (CORNISH, http://www.ast.leeds.ac.uk/Cornish/public/index.php) ('''B''' configuration), or the NRA VLA Sky Survey (NVSS, http://www.cv.nrao.edu/nvss/) ('''D''' configuration, all-sky) surveys. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 ([[#Documentation|Documentation]]) for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to obtain EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may find that such dissertations comprise pieces of several short proposals, which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being impaired by an adverse review of one proposal when the full scope of the project is not seen. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. Starting in 2011 time on the EVLA is scheduled on a semester basis, with each semester lasting six months. Proposal deadlines will be 5pm (1700) Eastern Time on February 1 and August 1 (if the deadline falls on a holiday or weekend, it is extended to the next working day). The February 1 proposal deadline nominally covers time to be scheduled during the following August through January, and the August 1 deadline is for time to be scheduled from February through July. Proposals for any configuration in the current '''D'''→'''A''' configuration cycle (September 2011 through January 2013) may be submitted at any proposal deadline, although a proposal for a configuration that has already passed may not be held over for consideration in the next configuration cycle, since the capabilities to be offered in the future are likely to be considerably different from those described in this document. All proposals will be reviewed by a Science Review Panel (SRP) in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The SRP's comments and rating are strongly advisory to the NRAO Time Allocation Committee (TAC), and the comments of both groups are passed on to the proposers soon after each meeting of the TAC (twice yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/observing/ for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2011 and 2012. == Director's Discretionary Time == The NRAO has established two categories of proposals for Director's Discretionary Time (DDT). DDT is limited to a maximum of 5% of the total observing time on the EVLA. All DDT proposals should be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Target of Opportunity.''' Target of Opportunity (ToO) proposals are for unexpected or unpredicted phenomena such as supernovae in nearby galaxies or extreme X-ray or radio flares. ToO Proposals are evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. ToO Proposals are evaluated on the basis of scientific merit by the Chair of the relevant Science Review Panel and Observatory staff with the necessary scientific expertise. The technical feasibility of the proposed observations will be assessed by Observatory staff. The proprietary period for data obtained by ToO Proposals will be assessed on a case-by-case basis but will be no more than six months. 2. '''Exploratory Time.''' Exploratory Proposals are normally for requests of small amounts of time, typically a few hours or less, in response to a recent discovery, possibly to facilitate future submission of a larger proposal. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current EVLA configuration rather than waiting 16 months. The possibility that a proposer forgets about or misses a proposal deadline, or just discovered that he/she was granted time for a particular source on some other telescope, will not constitute sufficient justification for granting observing time by this process. Thus, Exploratory Proposals must include a clear description of why the proposal could not have been submitted for normal review at a previous NRAO proposal deadline, and why it should not wait for the next proposal deadline. Proposals for exploratory time will be evaluated on the basis of scientific merit by the relevant Science Review Panel. Observatory staff will assess their technical feasibility. Notification of the disposition of an Exploratory Proposal normally will be within three weeks of receipt of the proposal; some of these proposals may be put in a queue such that they may or may not be observed. The proprietary period for data obtained by Exploratory Proposals normally will be six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of two formats: :– As a CASA Measurement Set. :– In UVFITS format, which can be read by either AIPS or CASA. The raw SDM format will only be available by special request. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. See http://casa.nrao.edu for more information on the latest release. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aips.nrao.edu for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/nsf06316/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Student Observing Support Program == In addition to travel support for individual data reduction visits NRAO maintains a program to support research by students, both graduate and undergraduate, at U.S. universities and colleges. Regular and Large proposals submitted for the EVLA, VLBA, and GBT, and any combination of these telescopes, are eligible. New applications to the program may be submitted along with new observing proposals at any proposal deadline. Details of this program can be found at https://science.nrao.edu/opportunities/student-programs/studentprograms == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aips.nrao.edu/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aips.nrao.edu/goaips.html. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa_cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS; data reduction and imaging algorithms |- | Miriam Krauss || 7230 || 300 || EVLA CASA subsystem scientist; rapid response science |- | Chris Langley || 7145 || 328 || EVLA Project Manager |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning; EVLA user support |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || WIDAR subsystem scientist; EVLA scientific software |- | Debra Shepherd || 7315 || 330 || EVLA Commissioning |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: queries should generally be directed to the NRAO Helpdesk, at http://science.nrao.edu/observing/helpdesk.shtml. However, you may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the editors of the present document (gvanmoor at nrao dot edu, cchandle at nrao dot edu) with questions on the material, or suggestions that would enhance the clarity of this guide. 3b1bfc687f468952263f116362e6a99a4c8138c6 1223 1222 2012-05-24T15:47:18Z Jott 8 wikitext text/x-wiki '''REPLACED ON 23 MAY 2012 BY https://science.nrao.edu/facilities/evla/docs/manuals/oss''' '''WE RECOMMEND YOU UPDATE YOUR BOOKMARKS ACCORDINGLY''' ---- '''The EVLA Observational Status Summary''' ''Version date: November 29, 2011'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget at the end of 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || align='center'| 10 μJy || align='center'| 1 μJy || align='center'| 10 |- | Maximum BW in each polarization || align='center'| 0.1 GHz || align='center'| 8 GHz || align='center'| 80 |- | Number of frequency channels at max. BW || align='center'| 16 || align='center'| 16,384 || align='center'| 1024 |- | Maximum number of freq. channels || align='center'| 512 || align='center'| 4,194,304 || align='center'| 8192 |- | Coarsest frequency resolution || align='center'| 50 MHz || align='center'| 2 MHz || align='center'| 25 |- | Finest frequency resolution || align='center'| 381 Hz || align='center'| 0.12 Hz || align='center'| 3180 |- | Number of full-polarization sub-correlators || align='center'| 2 || align='center'| 64 || align='center'| 32 |- | Log (Frequency Coverage over 1–50 GHz) || align='center'| 22% || align='center'| 100% || align='center'| 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date !Status or Actual Date |- | Installation of EVLA correlator subset for early science || align='center'| 2010 Q1 || align='center' | 2012 Q1 |- | Shared Risk Observing begins || align='center'| 2010 Q1 || align='center' | 2012 Q2 |- | Last antenna retrofitted || align='center'| 2010 Q2 || align='center' | 2010 Q2 |- | Full EVLA correlator installation || align='center'| 2011 Q2 || align='center' | 2011 Q2 |- | Last receiver installed || align='center'| 2012 Q4 || align='center' | on schedule |- |} == VLA to EVLA Transition == The year 2010 was extremely exciting for the EVLA. The correlator that was the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from '''A'''→'''B'''→'''C'''→'''D'''→'''A''' to '''D'''→'''C'''→'''B'''→'''A'''→'''D''', in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. The last VLA antenna was retrofitted to EVLA specifications in May 2010. During 2011 the WIDAR correlator was put into full observing mode with commissioning and Resident Shared Risk Observing starting in early 2011. By the end of 2011, up to 2 GHz of bandwidth was provided to the the general public along with 2 tunable bands, each with 8 spectral windows (64 to 256 channels) at S, Ku and X bands. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it has not yet been commissioned. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted '''D''', '''C''', '''B''', and '''A''' respectively. In addition, there are 3 “hybrid” configurations labelled '''DnC''', '''CnB''', and '''BnA''', in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle was modified during 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule for 2011, 2012 and 2013 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Approximate EVLA Configuration Schedule for 2011-2012''' ! Year ! Jan-Apr ! May-Aug ! Sep-Dec |- | align='center'| 2011 || align='center'| '''B''' || align='center'| '''A''' || align='center'| '''D''' |- | align='center'| 2012 || align='center'| '''C''' || align='center'| '''B''' || align='center'| '''A''' |- | align='center'| 2013 || align='center'| '''D''' || align='center'| '''C''' || align='center'| '''B''' |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 2012 (a full '''D'''→'''A''' configuration cycle) are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science is provided by two programs for outside users and one for EVLA commissioning staff. All early science programs are peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period thus involves an element of risk associated with the large stepwise increases in throughput bandwidth that are offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program provides early science capabilities to the general user community. These capabilities initially provided a maximum 256 MHz bandwidth that increased to 2 GHz for the '''D''' configuration in mid-2011 and will increase further to 8 GHz at the end of 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All EVLA antennas are outfitted with either EVLA or “interim” L, EVLA C, VLA X, EVLA K, EVLA Ka, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of July 2012, 25 of the EVLA antennas will be outfitted with S-band receivers, 23 EVLA antennas will have Ku-band receivers, and 22 will have new EVLA-style X-band receivers. Figure 1 shows the expected installation rate of final EVLA receiver systems for the rest of the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers, expected to be commissioned by the end of 2012. Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:RcvrAvailDec12.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of final EVLA receivers during the calendar year 2012 until the end of the EVLA Construction Project on December 31, 2012. Interim receivers with reduced frequency coverage or polarization purity are available at some bands in addition to those plotted (see Table 4).]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that are available in September 2011 and January 2012, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, Jan 2012 ! colspan="2" | Receiver availability, Jul 2012 |- ! ! GHz ! Final EVLA ! EVLA+VLA/interim ! Final EVLA ! EVLA+VLA/interim |- | 400 cm (4-band) || align='center'| 0.062–0.078 || align='center'| - || align='center'| - || align='center'|- || align='center'|- |- | 20 cm (L) || align='center'| 1.0–2.0 || align='center'| 19 || align='center'| 27 || align='center'|24 || align='center'|27 |- | 13 cm (S) || align='center'| 2.0–4.0 || align='center'| 21 || align='center'| 21 || align='center'|25 || align='center'|25 |- | 6 cm (C) || align='center'| 4.0–8.0 || align='center'| 27 || align='center'| 27 || align='center'|27 || align='center'|27 |- | 3 cm (X) || align='center'| 8.0–12.0 || align='center'| 16 || align='center'| 27 || align='center'|22 || align='center'|27 |- | 2 cm (Ku) || align='center'| 12.0–18.0 || align='center'| 18 || align='center'| 18 || align='center'|23 || align='center'|23 |- | 1.3 cm (K) || align='center'| 18.0–26.5 || align='center'| 27 || align='center'| 27 || align='center'|27 || align='center'|27 |- | 1 cm (Ka) || align='center'| 26.5–40.0 || align='center'| 27 || align='center'| 27 || align='center'|27 || align='center'|27 |- | 0.7 cm (Q) || align='center'| 40.0–50.0 || align='center'| 27 || align='center'| 27 || align='center'|27 || align='center'|27 |- |} :Note: The "EVLA+VLA/interim" columns give the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers (see Figure 1). == Open Shared Risk Observing (OSRO) == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program has extended this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program obtain good quality data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing (RSRO) == The WIDAR correlator and the EVLA provides a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program is now expected to run through the end of 2012, with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. Full operations of the EVLA will begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program for the upcoming '''D'''→'''A''' configuration cycle, Sep 2011 through Dec 2012. Users interested in participating in the RSRO program should refer to the web page at https://science.nrao.edu/facilities/evla/early-science/rsro for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24 || 80 || 260 || 850 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka) || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800 || 2200 || 20000 || 20000 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka) || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations ('''DnC''', '''CnB''', '''BnA''') should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., '''DnB''', or '''CnA''') is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. For the '''C''' configuration an antenna from the middle of the north arm is moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Note that although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration. ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are insufficient EVLA 8–12 GHz receivers available yet to determine system performance across the band. A project with the goal of doubling the longest baseline available in the '''A''' configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the ''SEFD'' at some fiducial EVLA frequencies. [[File:SEFD.png|none|frame|Figure 2: ''SEFD'' for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the Ku, K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: ''SEFD''s and '''D'''-Configuration Confusion Limits''' !Frequency !''SEFD'' !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || align='center'| 420 || align='center'| 89 |- | 3.0 GHz (S) || align='center'| 370 || align='center'| 14 |- | 6.0 GHz (C) || align='center'| 310 || align='center'| 2.3 |- | 10.0 GHz (X) || align='center'| 250 || align='center'| negligible |- | 15 GHz (Ku) || align='center'| 350 || align='center'| negligible |- | 22 GHz (K) || align='center'| 560 || align='center'| negligible |- | 33 GHz (Ka) || align='center'| 730 || align='center'| negligible |- | 45 GHz (Q) || align='center'| 1400 || align='center'| negligible |} :Note: ''SEFD''s at Ku, K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in '''D''' configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temperature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted above assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. '''At source elevations greater than 80 degrees (zenith angle < 10 degrees), source tracking becomes difficult; it is recommended to avoid such source elevations during the observation preparation setup'''. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in Right Ascension (RA). The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where ''T''_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and ''S'' in mJy per beam, the constant ''F'' depends only upon array configuration and has the approximate value ''F'' = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for ''S''. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of ''F'' calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where ''D'' is the distance to the galaxy in Mpc, and ''S∆V'' is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different “baseband” frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” Currently, a maximum of 1.024 GHz can be correlated for each IF pair (see [[#Correlator Configurations|Correlator Configurations]]), for a total maximum bandwidth of approximately 2 GHz. The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. At X-band a number of the antennas will continue to have the old narrow-band VLA X-band receivers until their retrofit is complete at the end of the EVLA construction project. As of December 2010 there is not a sufficient number of EVLA-style X-band receivers in the array to evaluate either the system performance or the radio frequency interference environment throughout the EVLA X-band tuning range of 8-12 GHz. A total bandwidth of 800 MHz equivalent to that of the VLA receivers (8.0-8.8 GHz) should be assumed for the purposes of sensitivity calculations at X-band for the present. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || align='center'| 1.0–2.0 || align='center'| 1.25¹ || align='center'| 1.75¹ |- | 13 cm (S) || align='center'| 2.0–4.0 || align='center'| 2.5 || align='center'| 3.5 |- | 6 cm (C) || align='center'| 4.0–8.0 || align='center'| 5.0 || align='center'| 6.0 |- | 3 cm (X) || align='center'| 8.0–12.0 || align='center'| 8.5² || align='center'| 9.5² |- | 2 cm (Ku) || align='center'| 12.0–18.0 || align='center'| 13.5 || align='center'| 14.5 |- | 1.3 cm (K) || align='center'| 18.0–26.5 || align='center'| 20.7 || align='center'| 21.7 |- | 1 cm (Ka) || align='center'| 26.5–40.0 || align='center'| 31.5 || align='center'| 32.5 |- | 0.7 cm (Q) || align='center'| 40.0–50.0 || align='center'| 40.5 || align='center'| 41.5 |} :Notes: :: 1. This default frequency set-up for L-band comprises two 512 MHz basebands (each with 8 subbands of 64 MHz) to cover the entire 1-2 GHz of the L-band receiver. :: 2. Many of the antennas continue to have the old narrow-band VLA receivers (see Figure 1), for which a total bandwidth of 800 MHz should be assumed (8.0-8.8 GHz). The RFI environment of the default tunings has not yet been evaluated. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 in [[#Documentation|Documentation]] for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! (∆ν/ν₀)*(θ₀/θ_HPBW) ! Peak ! Width |- | 0.0 || align='center'| 1.00 || align='center'| 1.00 |- | 0.50 || align='center'| 0.95 || align='center'| 1.05 |- | 0.75 || align='center'| 0.90 || align='center'| 1.11 |- | 1.0 || align='center'| 0.80 || align='center'| 1.25 |- | 2.0 || align='center'| 0.50 || align='center'| 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 (in [[#Documentation|Documentation]]) considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || align='center'| 2.1 || align='center'| 4.8 || align='center'| 6.7 |- | '''B''' || align='center'| 6.8 || align='center'| 15.0 || align='center'| 21.0 |- | '''C''' || align='center'| 21.0 || align='center'| 48.0 || align='center'| 67.0 |- | '''D''' || align='center'| 68.0 || align='center'| 150.0 || align='center'| 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 (http://www.aoc.nrao.edu/evla/geninfo/memoseries/evlamemo67.pdf) for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second, for '''A''' configuration. For '''B''', '''C''', and '''D''' configurations the minimum integration time is 3 seconds, unless a shorter integration time is explicitly requested and justified. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 14.4 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 51.8 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 1.2 TB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool (https://science.nrao.edu/facilities/evla/data-archive/evla) will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on disk drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the '''D''' configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum at L-band. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots and tables of known RFI are available online, at http://science.nrao.edu/evla/observing/RFI/; plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the '''D''' configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in '''A''' configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in '''A''' configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List (http://www.vla.nrao.edu/astro/calib/manual/) should be used. The positions of these sources are taken from lists published by the United States Naval Observatory (USNO). == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1 in [[#Documentation|Documentation]]. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth smearing and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its '''D''' configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP (Wilkinson Microwave Anisotropy Probe) flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby calibrator with an "S" code in the calibrator database, and a more distant calibrator with a "P" code, the nearby calibrator is usually the better choice (see http://www.vla.nrao.edu/astro/calib/manual/key.html for a description of calibrator codes). :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173 (http://www.vla.nrao.edu/memos/sci/). These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API may be found at http://www.vla.nrao.edu/astro/guides/api/. Plots of current/historical data can be found at: https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi Characteristic seasonal averages are represented below: {| border="1" align="center" |+ '''Table: Seasonal API/wind values at the EVLA''' !Month !API (night) [deg] !API (median) [deg] !API (day) [deg] !Wind (night) [m/s] !Wind (median) [m/s] !Wind (day) [m/s] |- | [[Media:APIwind_January.png| January]] || 2.3 || 2.8 || 3.6 || 1.6 || 1.9 || 2.3 |- | [[Media:APIwind_February.png| February]] || 2.9 || 3.4 || 4.5 || 4.0 || 4.3 || 4.5 |- | [[Media:APIwind_March.png| March]] || 2.8 || 3.7 || 5.5 || 3.4 || 3.9 || 4.7 |- | [[Media:APIwind_April.png| April]] || 3.3 || 4.5 || 6.2 || 5.3 || 5.5 || 5.8 |- | [[Media:APIwind_May.png | May]] || 2.9 || 4.6 || 6.7 || 2.6 || 3.2 || 3.7 |- | [[Media:APIwind_June.png| June]] || 3.8 || 5.5 || 7.4 || 2.5 || 3.9 || 6.3 |- | [[Media:APIwind_July.png| July]] || 6.2 || 8.3 || 10.5 || 2.9 || 2.9 || 3.0 |- | [[Media:APIwind_August.png| August]] || 5.4 || 7.1 || 11.3 || 1.7 || 2.3 || 3.0 |- | [[Media:APIwind_September.png| September]] || 5.2 || 6.6 || 8.8 || 2.3 || 3.0 || 3.6 |- | [[Media:APIwind_October.png| October]] || 4.2 || 5.3 || 7.4 || 2.3 || 2.9 || 3.7 |- | [[Media:APIwind_November.png| November]] || 2.6 || 3.0 || 4.0 || 1.2 || 2.5 || 1.6 |- | [[Media:APIwind_December.png| December]] || 2.8 || 3.2 || 4.1 || 1.2 || 1.6 || 2.7 |- |} Click on the Month links above to see plots of phase and wind speed versus time. :Note: day indicates sunrise to sunset values; night indicates sunset to sunrise values. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/<math>{\rm{m}^2}</math>, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. For more details see: https://science.nrao.edu/facilities/evla/early-science/polarimetry == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. For the Open Shared Risk Observing (OSRO) program available to the community during the period Sep 2011 through Jan 2013 are offering two independently tunable basebands, where each baseband has up to eight sub-bands. Possible sub-band widths are 128 MHz, 64 MHz, 32 MHz, all the way down in factors of 2 to 0.03125 MHz. All sub-bands must have the same bandwidth and channelization in both basebands, and be contiguous in frequency within each baseband. We will offer three different OSRO modes: full polarization, dual polarization, and single polarization, with 64, 128, and 256 channels per sub-band, respectively. There is always the possibility during offline processing to smooth in frequency to reduce dataset sizes or to improve spectral response. Starting with the '''D'''-configuration in September 2011, we have been providing options for configuring WIDAR for OSRO in the following three ways: :1. “OSRO Full Polarization”: Four polarization products. This configuration offers 4 polarization products for each sub-band, each of which has 128 MHz bandwidth with 64 channels. It is possible to decrease the subband bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in the following table: {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO Full Polarization)''' ! Sub-band BW (MHz) ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 64 || 1000 || 300 || 19,200 |- | 32 || 64 || 500 || 150 || 9,600 |- | 16 || 64 || 250 || 75 || 4,800 |- | 8 || 64 || 125 || 37.5 || 2,400 |- | 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 64 || 31.25 || 9.4 || 600 |- | 1 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO Dual Polarization”: Two polarization products. This configuration offers 2 polarization products for each sub-band, each of which has 128 MHz bandwidth with 128 channels. It is possible to decrease the sub-band bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in the following table. {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for dual polarization (OSRO Dual Polarization)''' ! Sub-band BW (MHz) ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 128 || 1000 || 300/ν (GHz) || 38,400/ν (GHz) |- | 64 || 128 || 500 || 150 || 19,200 |- | 32 || 128 || 250 || 75 || 9,600 |- | 16 || 128 || 125 || 37.5 || 4,800 |- | 8 || 128 || 62.5 || 19 || 2,400 |- | 4 || 128 || 31.25 || 9.4 || 1,200 |- | 2 || 128 || 15.625 || 4.7 || 600 |- | 1 || 128 || 7.813 || 2.3 || 300 |- | 0.5 || 128 || 3.906 || 1.2 || 150 |- | 0.25 || 128 || 1.953 || 0.59 || 75 |- | 0.125 || 128 || 0.977 || 0.29 || 37.5 |- | 0.0625 || 128 || 0.488 || 0.15 || 18.75 |- | 0.03125 || 128 || 0.244 || 0.073 || 9.375 |} :3: "OSRO Single Polarization": One polarization product (new for OSRO observing). It offers 1 polarization product for each sub-band, each of which has 128 MHz bandwidth with 256 channels. It is possible to decrease the sub-band bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in the following table. {| border="1" align="center" |+ '''Table 14: Correlator capabilities per sub-band for single polarization (OSRO Single Polarization)''' ! Sub-band BW (MHz) ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 256 || 250 || 75 || 19,200 |- | 32 || 256 || 125 || 37.5 || 9,600 |- | 16 || 256 || 62.5 || 19 || 4,800 |- | 8 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities are being provided with integration times no shorter than 1 second in '''A''' configuration (3 seconds in '''B'''/'''C'''/'''D''' configurations), and Doppler setting will be available with these correlator configurations. If it is likely that the data will need to be resampled spectrally in order to Doppler track to a line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna (e.g., Y1) modes, have not yet been commissioned and are not yet available to the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing "Faint Images of the Radio Sky at Twenty-centimeters survey(FIRST, http://www.cv.nrao.edu/first/) or the Co-Ordinated Radio 'N' Infrared Survey for High-mass star formation (CORNISH, http://www.ast.leeds.ac.uk/Cornish/public/index.php) ('''B''' configuration), or the NRA VLA Sky Survey (NVSS, http://www.cv.nrao.edu/nvss/) ('''D''' configuration, all-sky) surveys. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 ([[#Documentation|Documentation]]) for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to obtain EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may find that such dissertations comprise pieces of several short proposals, which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being impaired by an adverse review of one proposal when the full scope of the project is not seen. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. Starting in 2011 time on the EVLA is scheduled on a semester basis, with each semester lasting six months. Proposal deadlines will be 5pm (1700) Eastern Time on February 1 and August 1 (if the deadline falls on a holiday or weekend, it is extended to the next working day). The February 1 proposal deadline nominally covers time to be scheduled during the following August through January, and the August 1 deadline is for time to be scheduled from February through July. Proposals for any configuration in the current '''D'''→'''A''' configuration cycle (September 2011 through January 2013) may be submitted at any proposal deadline, although a proposal for a configuration that has already passed may not be held over for consideration in the next configuration cycle, since the capabilities to be offered in the future are likely to be considerably different from those described in this document. All proposals will be reviewed by a Science Review Panel (SRP) in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The SRP's comments and rating are strongly advisory to the NRAO Time Allocation Committee (TAC), and the comments of both groups are passed on to the proposers soon after each meeting of the TAC (twice yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/observing/ for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2011 and 2012. == Director's Discretionary Time == The NRAO has established two categories of proposals for Director's Discretionary Time (DDT). DDT is limited to a maximum of 5% of the total observing time on the EVLA. All DDT proposals should be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Target of Opportunity.''' Target of Opportunity (ToO) proposals are for unexpected or unpredicted phenomena such as supernovae in nearby galaxies or extreme X-ray or radio flares. ToO Proposals are evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. ToO Proposals are evaluated on the basis of scientific merit by the Chair of the relevant Science Review Panel and Observatory staff with the necessary scientific expertise. The technical feasibility of the proposed observations will be assessed by Observatory staff. The proprietary period for data obtained by ToO Proposals will be assessed on a case-by-case basis but will be no more than six months. 2. '''Exploratory Time.''' Exploratory Proposals are normally for requests of small amounts of time, typically a few hours or less, in response to a recent discovery, possibly to facilitate future submission of a larger proposal. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current EVLA configuration rather than waiting 16 months. The possibility that a proposer forgets about or misses a proposal deadline, or just discovered that he/she was granted time for a particular source on some other telescope, will not constitute sufficient justification for granting observing time by this process. Thus, Exploratory Proposals must include a clear description of why the proposal could not have been submitted for normal review at a previous NRAO proposal deadline, and why it should not wait for the next proposal deadline. Proposals for exploratory time will be evaluated on the basis of scientific merit by the relevant Science Review Panel. Observatory staff will assess their technical feasibility. Notification of the disposition of an Exploratory Proposal normally will be within three weeks of receipt of the proposal; some of these proposals may be put in a queue such that they may or may not be observed. The proprietary period for data obtained by Exploratory Proposals normally will be six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of two formats: :– As a CASA Measurement Set. :– In UVFITS format, which can be read by either AIPS or CASA. The raw SDM format will only be available by special request. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. See http://casa.nrao.edu for more information on the latest release. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aips.nrao.edu for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/nsf06316/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Student Observing Support Program == In addition to travel support for individual data reduction visits NRAO maintains a program to support research by students, both graduate and undergraduate, at U.S. universities and colleges. Regular and Large proposals submitted for the EVLA, VLBA, and GBT, and any combination of these telescopes, are eligible. New applications to the program may be submitted along with new observing proposals at any proposal deadline. Details of this program can be found at https://science.nrao.edu/opportunities/student-programs/studentprograms == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aips.nrao.edu/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aips.nrao.edu/goaips.html. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa_cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS; data reduction and imaging algorithms |- | Miriam Krauss || 7230 || 300 || EVLA CASA subsystem scientist; rapid response science |- | Chris Langley || 7145 || 328 || EVLA Project Manager |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning; EVLA user support |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || WIDAR subsystem scientist; EVLA scientific software |- | Debra Shepherd || 7315 || 330 || EVLA Commissioning |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: queries should generally be directed to the NRAO Helpdesk, at http://science.nrao.edu/observing/helpdesk.shtml. However, you may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the editors of the present document (gvanmoor at nrao dot edu, cchandle at nrao dot edu) with questions on the material, or suggestions that would enhance the clarity of this guide. 6cd46c35ddca47a13e2a9e5ecedfb039c06f6bcc 10 2 1225 1032 2012-05-24T15:49:38Z Jott 8 wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] Key Links<br> [http://science.nrao.edu/evla/ '''EVLA web-site'''] · [https://staff.nrao.edu/wiki/bin/view/EVLA/EVLACommissioning '''EVLA Commissioning and Science Verification Twiki'''] |- |valign=top|[[Image:book.gif]] EVLA Information<br> [https://science.nrao.edu/facilities/evla/docs/manuals/oss '''Observational Status Summary'''] · [http://www.aoc.nrao.edu/cgi-bin/evla/receivers-report '''EVLA Antenna-Receiver Availability'''] · [[EVLA Acronym List]] |- |valign=top|[[Image:book.gif]] EVLA Utilities<br> [https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi '''EVLA Atmospheric Phase/Wind Monitoring'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotpointing.cgi '''EVLA Pointing solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/plotgains.cgi '''EVLA Gain solution plotter'''] · [https://webtest.aoc.nrao.edu/cgi-bin/thunter/bptool.cgi '''EVLA Bandpass plotter'''] |- |valign=top|[[Image:book.gif]] EVLA Observing Preparation<br> [[:Category:SpectraLine| Spectral Line Observations]] · [[:Category:Polarimetry| Polarimetry Observations]] · [[:Category:Planetary| Planetary Observations]] · [[:Category:HighFrequency| High Frequency Observing (K, Ka, Q)]] · [[:Category:LowFrequency| Low Frequency Observing (L, S, C)]] · [[:Category:PhasedArray| Phased Array Observing]] [[:Category:Pulsar| Pulsar Observing]] · [[:Category:OPT-QuickStart| OPT Quick Start Guide]] |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> CASA: [http://casaguides.nrao.edu '''CASA Reduction Guides'''] <br> AIPS: [[Key to Calcodes]] · [[:Category:Post-Processing|Special Considerations for EVLA Data Calibration and Imaging in AIPS]]<br> |} d20e3d8a4a2c5eaacea5f4366b331131baa558b1 8 126 1238 2012-06-13T14:26:25Z Admin 1 Created page with "/* CSS placed here will be applied to GeSHi syntax highlighting */ div.mw-geshi { background-color: #ffe4b5; }" css text/css /* CSS placed here will be applied to GeSHi syntax highlighting */ div.mw-geshi { background-color: #ffe4b5; } 5adb7bbf0c18447e07fbff2d3f27b5061052a851 6 127 1246 2012-06-20T06:46:56Z Jott 8 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 1249 1246 2012-06-20T06:50:43Z Jott 8 Jott uploaded a new version of &quot;[[File:ExposureCalculator.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 128 1254 2012-06-20T07:15:24Z Jott 8 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 129 1255 2012-06-20T07:15:46Z Jott 8 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 130 1263 2012-06-20T08:06:46Z Jott 8 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 1264 1263 2012-06-20T08:08:55Z Jott 8 Jott uploaded a new version of &quot;[[File:Doppler.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 0 48 1269 1268 2012-06-27T05:08:25Z Jott 8 /* Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = cz </math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 39cc198e00eacb58d370883602c3e7bba4319d51 1270 1269 2012-06-27T05:08:42Z Jott 8 /* Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> <math> v^{radio} = \frac{\lambda_0-\nu}{\nu_0}\,\,c = cz </math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> acc27a6fcb401d2c76ddb1a84fa6380077fabd37 1271 1270 2012-06-27T05:09:01Z Jott 8 /* Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> <math> v^{optical} = \frac{\lambda_0-\nu}{\nu_0}\,\,c = cz </math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 5d1a004e6f41c7d9be0111fa4f06ddd366356c72 1272 1271 2012-06-27T05:09:25Z Jott 8 /* Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> cab26c8da0a2b144b4f382a5c7d1ce2b9f861fe7 1273 1272 2012-06-27T05:09:51Z Jott 8 /* Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> <math> v^{optical} = \frac{\nu_0-\lambda}{\lambda_0}\,\,c = cz </math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 8fc6487e3eb5758941e0124e40a78a7c68801837 1274 1273 2012-06-27T05:11:04Z Jott 8 /* Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> <math> v^{optical} = \frac{\nu_{0}-\lambda}{\lambda_0}\,\,c = cz </math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 010abbc9f5d2f058e1b80b2470e733c4202a848c 1275 1274 2012-06-27T05:11:24Z Jott 8 /* Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> <math> v^{optical} = \frac{\lambda_{0}-\lambda}{\lambda_0}\,\,c = cz </math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 4beed304a05c5cff311a7dacc557eb3fe7a5990b 1276 1275 2012-06-27T05:11:49Z Jott 8 /* Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> <math> v^{optical} = \frac{\lambda_0-\lambda}{\lambda_0}\,\,c = cz </math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> cab26c8da0a2b144b4f382a5c7d1ce2b9f861fe7 1277 1276 2012-06-27T05:13:49Z Jott 8 /* Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth oribtal motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> c0662672ff8bcd61ffc543d1c4f04064f0c52f37 1278 1277 2012-06-27T05:16:30Z Jott 8 /* Velocity Frames */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g. At a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the VLA. <i>The VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 8fb05c4b2d412c33d1e8603ec3e4b22fd6a56437 1279 1278 2012-06-27T05:21:07Z Jott 8 /* Doppler Correction */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weigthed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoohting was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task <i>hanningsmooth</i>. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 4bec520d11a9799c990088a87245aca1e8535f14 1280 1279 2012-06-27T05:22:47Z Jott 8 /* Gibbs Phenomenon and Hanning smoothing */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should be the correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We like to refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advise on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> d581bcb43bb47c6eb0ad3b7b50dc882b6c8b7f7b 1281 1280 2012-06-27T05:24:02Z Jott 8 /* Sensitivity/Exposure Time Calculation */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 254c5e0ecdfaad201977980425afc951e972147b 1282 1281 2012-06-27T05:25:01Z Jott 8 /* Basebands */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128 MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 9535bcc9534c9530cf4939ac7a2aedb2afa5da57 1283 1282 2012-06-27T05:25:39Z Jott 8 /* Fixed 128MHz Subbands and 128 MHz "Suckouts" */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures. To obtain more spectral channels, one can consider to use single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 239d3c8d19796d5336f0fa63bde4f650d69b2d2f 1284 1283 2012-06-27T05:28:38Z Jott 8 /* Fixed 128MHz Subbands and 128MHz "Suckouts" */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). Fundamentally, each narrow subband can be between 128 MHz and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channel in dual polariation mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables below display the options and the corresponding velocity widths and channelizations: * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHx fixed subband. '''(??? Michael to comment - what other restrictions are there ???)''' === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> ab157b79faffb736ea031d5e2c23951c74497fa2 1285 1284 2012-06-27T05:37:47Z Jott 8 /* Narrow Subbands */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> c11f787f8d413e6094ab79e8eca1f1bdeb1fd29f 1286 1285 2012-06-27T05:38:57Z Jott 8 /* Correlator Baselineboards */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "Bl.BPs" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 480a9e8326b0cf5b9ed1625eb1e55ae844e910bd 1287 1286 2012-06-27T05:39:46Z Jott 8 /* Correlator Baselineboards */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "Bl.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> fbbeb2049b66f48715c5ac551644cec753b0d0b0 1288 1287 2012-06-27T05:40:40Z Jott 8 /* Correlator Baselineboards */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BlB.BPS dropdown menu that is available for each subband. Baselineboard stacking is extremely useful as it allows to use wider bandwidths in each subband but yet maintain a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. Alternatively, for a given subband bandwidth, BlB stacking provides a higher spectral resolution. E.g. the full 2GHz bandwidth of the 8bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as of the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reducies the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> d992bc1589d460060925d5868ad7071bb712337f 1289 1288 2012-06-27T05:46:55Z Jott 8 /* Baselineboard Stacking */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available in for shared risk observations. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> f07c84f3c2069b60ade8bb72cc0455c9daa54223 1290 1289 2012-06-27T05:48:01Z Jott 8 /* Recirculation */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands an polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To observe all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 5aa1885e952f21f1b4aa21418ccb6a8f3317e996 1291 1290 2012-06-27T05:51:00Z Jott 8 /* Planning and Setup of VLA Spectral Line Observations */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== Reminder: For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * up to 16 subband pairs per baseband pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be setup with the flexibility mentioned above. The figure shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes.''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only hve access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 75 MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 045c06c115871867f4e4ade2e4150e5197af2332 1292 1291 2012-06-27T06:01:07Z Jott 8 /* Planning and Setup of VLA Spectral Line Observations */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 65 MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> e6703cf722eead62789873a13fde70207d78dbde 1293 1292 2012-06-27T06:01:42Z Jott 8 /* Data Rate Limits */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available. During commissioning, it is advisable to not use all baseline boards, as some of them may undergo maintenance. Also make sure to stay within the data rate restriction limits. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In teh current OSRO/RSRO interface, the subbands are not independently tuneable yet (but this feature will be available for the Ausgust 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should afll in that subband. Those entries are not used anywhere outside OPT '''(??? rest frequencies for OSRO ???)'''. The ''delete'' button removes the subband if is is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as the script currently need manual editing after OPT submission''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> b37d65188e681f1f24e0ac40ac9a3420ae27a61d 1294 1293 2012-06-27T06:07:04Z Jott 8 /* Frequency Setup */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. Contradictory to to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 9d597ff1e40597dd41d12ab2981e8c1a5e1efbf6 1295 1294 2012-06-27T06:10:09Z Jott 8 /* Doppler Setting */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies, however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to the per cent level. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> a55b0cbad9b744674024a5b5127b40aa0203e288 1296 1295 2012-06-27T06:26:07Z Jott 8 /* Bandpass Setup */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on this [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 84cfe5ae1bf9617dc33fdfee3182b75042baf930 1297 1296 2012-06-27T06:26:45Z Jott 8 /* Bandpass Setup */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 07dfb4d2f38880d381bc46bd9ab7ea6c0b541120 1298 1297 2012-06-27T06:39:29Z Jott 8 /* Continuum Subtraction */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution such that the ringing effects beat against each other, effectively reduce the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> fd1f0ab2f2b11fe5bc8f781ab223770cbe0bf67e 1299 1298 2012-06-27T06:42:01Z Jott 8 /* Continuum Subtraction */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> d4c5f94e30ff0ffe358fb97af6e5e5731a00ebbb 1300 1299 2012-07-02T04:12:40Z Jott 8 wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == 1 Introducing VLA Spectral Line Observing == The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing <!-- JUERGEN -indent this --> two 1GHz basebands <!-- JUERGEN -indent this --> up to 16 independently tunable subbands per baseband <!-- JUERGEN -indent this --> independent subband bandwidths ranging from 31.5kHz to 128MHz <!-- JUERGEN -should this be 31.25kHz-128MHz? --> <!-- JUERGEN -indent this --> independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- <!-- JUERGEN - I did 2048 channels in dual pol for my RSRO. Is this possible for regular observing for Aug 2012 deadline? --> --> * 3-bit samplers providing <!-- JUERGEN -indent this --> four 2GHz basebands with dual polarization, 128MHz subbands, and 2MHz resolution <!--JUERGEN -is this right? --> == 2 Line Frequencies and Velocity Conventions == === 2.1 Line Rest Frequencies === Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == 2.2 Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> <!-- JUERGEN -make sure this shows up as a v, not a nu -->) and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift.<!--- JUERGEN - rewrite this sentence, its jumbled and unclear. --> == 2.3 Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Therefore in order to determine the observing frequency, one must specify the frame in which the source velocity is measured. Various rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] <!-- JUERGEN - why the extra ]'s?--> == 2.4 Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accomodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == 3 The WIDAR Correlator == === 3.1 Basebands === Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independantly tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] ==== 3.1.1 Baseband Tuning Restrictions ==== The following restrictions apply to baseband tuning: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * With 3-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD <!--JUERGEN -which BD? B1D1 or B2D2? --> * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation <!-- JUERGEN - 8-bit samplers should have AC0, not A1C1, B1D1, B2D2, etc. --> === 3.2 Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands, which actually measure the spectrum, can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. The preferred method is to avoid them in your spectral setup. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === 3.3 Correlator Resources and Subband Placement === Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 BlB pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. ==== 3.3.1 Narrow Subbands with the 8-bit sampler ==== Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This means that ALL subbands, All narrow subbands can overlap each other, and every subband is independent, so that <i> different subbands may have different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> <!-- JUERGEN - indent this sectoin maybe. --> ===== Standard Subbands ===== Standard subbands allocate a single baselineboard pair to a each subband (in single polarizations) or subband pair (in dual or full polarization). Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 <!--JUERGEN 64 total subband pairs for shared risk yes? -->) for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> <!-- JUERGEN - indent this sectoin maybe. --> ===== Baselineboard Stacking ===== In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baselineboards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline, however we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> <!-- JUERGEN - indent this sectoin maybe. --> ===== Recirculation ===== "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. ==== 3.3.1 Narrow Subbands with the 3-bit sampler ==== As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs with dual polarization, 2MHz resolution and 128MHz subbands. <!-- JUERGEN -anything else to say? Is this only dual polarization? --> === 3.4 Data Rate Limits === Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. == 4 Tips for Planning, Setup, and Processing of VLA Spectral Line Observations == '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing <!-- JUERGEN -indent this --> two 1GHz baseband pairs <!-- JUERGEN -indent this --> up to 16 independently tunable subband pairs per baseband pair <!-- JUERGEN -indent this --> independent subband bandwidths ranging from 31.25kHz to 128MHz <!-- JUERGEN -should this be 31.25kHz-128MHz? --> <!-- JUERGEN -indent this --> independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN -once again, I did 2048 channels in dual pol. Is this possible? --> * 3-bit samplers providing <!-- JUERGEN -indent this --> four 2GHz basebands with dual polarization, 128MHz subband bandwidths, and 2MHz resolution <!--JUERGEN -is this right? --> === 4.1 Considerations for Planning Subband Bandwidths and Resolution ==== 4.1.1 Bandwidths required for UV continuum subtraction ==== When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. ==== 4.1.2 Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. === 4.2 Sensitivity/Exposure Time Calculation === [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. === 4.3 The Proposal Submission Tool (PST) === The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === 4.4 Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. ==== 4.4.1 Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> ==== 4.4.2 Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== 4.4.3 Correlator Resources Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 44b6f6d999cdf292341c87d176c3f4784980f23d 1301 1300 2012-07-02T04:19:27Z Jott 8 /* 1 Introducing VLA Spectral Line Observing */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == 1 Introducing VLA Spectral Line Observing == The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN - I did 2048 channels in dual pol for my RSRO. Is this possible for regular observing for Aug 2012 deadline? --> * 3-bit samplers providing ** four 2GHz basebands with dual polarization, 128MHz subbands, and 2MHz resolution <!--JUERGEN -is this right? --> == 2 Line Frequencies and Velocity Conventions == === 2.1 Line Rest Frequencies === Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == 2.2 Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> <!-- JUERGEN -make sure this shows up as a v, not a nu -->) and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift.<!--- JUERGEN - rewrite this sentence, its jumbled and unclear. --> == 2.3 Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Therefore in order to determine the observing frequency, one must specify the frame in which the source velocity is measured. Various rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] <!-- JUERGEN - why the extra ]'s?--> == 2.4 Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accomodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == 3 The WIDAR Correlator == === 3.1 Basebands === Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independantly tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] ==== 3.1.1 Baseband Tuning Restrictions ==== The following restrictions apply to baseband tuning: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * With 3-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD <!--JUERGEN -which BD? B1D1 or B2D2? --> * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation <!-- JUERGEN - 8-bit samplers should have AC0, not A1C1, B1D1, B2D2, etc. --> === 3.2 Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands, which actually measure the spectrum, can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. The preferred method is to avoid them in your spectral setup. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === 3.3 Correlator Resources and Subband Placement === Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 BlB pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. ==== 3.3.1 Narrow Subbands with the 8-bit sampler ==== Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This means that ALL subbands, All narrow subbands can overlap each other, and every subband is independent, so that <i> different subbands may have different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> <!-- JUERGEN - indent this sectoin maybe. --> ===== Standard Subbands ===== Standard subbands allocate a single baselineboard pair to a each subband (in single polarizations) or subband pair (in dual or full polarization). Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 <!--JUERGEN 64 total subband pairs for shared risk yes? -->) for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> <!-- JUERGEN - indent this sectoin maybe. --> ===== Baselineboard Stacking ===== In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baselineboards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline, however we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> <!-- JUERGEN - indent this sectoin maybe. --> ===== Recirculation ===== "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. ==== 3.3.1 Narrow Subbands with the 3-bit sampler ==== As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs with dual polarization, 2MHz resolution and 128MHz subbands. <!-- JUERGEN -anything else to say? Is this only dual polarization? --> === 3.4 Data Rate Limits === Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. == 4 Tips for Planning, Setup, and Processing of VLA Spectral Line Observations == '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing <!-- JUERGEN -indent this --> two 1GHz baseband pairs <!-- JUERGEN -indent this --> up to 16 independently tunable subband pairs per baseband pair <!-- JUERGEN -indent this --> independent subband bandwidths ranging from 31.25kHz to 128MHz <!-- JUERGEN -should this be 31.25kHz-128MHz? --> <!-- JUERGEN -indent this --> independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN -once again, I did 2048 channels in dual pol. Is this possible? --> * 3-bit samplers providing <!-- JUERGEN -indent this --> four 2GHz basebands with dual polarization, 128MHz subband bandwidths, and 2MHz resolution <!--JUERGEN -is this right? --> === 4.1 Considerations for Planning Subband Bandwidths and Resolution ==== 4.1.1 Bandwidths required for UV continuum subtraction ==== When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. ==== 4.1.2 Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. === 4.2 Sensitivity/Exposure Time Calculation === [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. === 4.3 The Proposal Submission Tool (PST) === The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === 4.4 Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. ==== 4.4.1 Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> ==== 4.4.2 Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== 4.4.3 Correlator Resources Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> c8eaa64da5a6f033e0dd06caaf4db9bbf5563cb7 1302 1301 2012-07-02T04:36:37Z Jott 8 /* 2.2 Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == 1 Introducing VLA Spectral Line Observing == The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN - I did 2048 channels in dual pol for my RSRO. Is this possible for regular observing for Aug 2012 deadline? --> * 3-bit samplers providing ** four 2GHz basebands with dual polarization, 128MHz subbands, and 2MHz resolution <!--JUERGEN -is this right? --> == 2 Line Frequencies and Velocity Conventions == === 2.1 Line Rest Frequencies === Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == 2.2 Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift to the channel and line widths as required for the redshifted line. == 2.3 Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Therefore in order to determine the observing frequency, one must specify the frame in which the source velocity is measured. Various rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] <!-- JUERGEN - why the extra ]'s?--> == 2.4 Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accomodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == 3 The WIDAR Correlator == === 3.1 Basebands === Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independantly tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] ==== 3.1.1 Baseband Tuning Restrictions ==== The following restrictions apply to baseband tuning: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * With 3-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD <!--JUERGEN -which BD? B1D1 or B2D2? --> * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation <!-- JUERGEN - 8-bit samplers should have AC0, not A1C1, B1D1, B2D2, etc. --> === 3.2 Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands, which actually measure the spectrum, can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. The preferred method is to avoid them in your spectral setup. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === 3.3 Correlator Resources and Subband Placement === Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 BlB pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. ==== 3.3.1 Narrow Subbands with the 8-bit sampler ==== Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This means that ALL subbands, All narrow subbands can overlap each other, and every subband is independent, so that <i> different subbands may have different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> <!-- JUERGEN - indent this sectoin maybe. --> ===== Standard Subbands ===== Standard subbands allocate a single baselineboard pair to a each subband (in single polarizations) or subband pair (in dual or full polarization). Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 <!--JUERGEN 64 total subband pairs for shared risk yes? -->) for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> <!-- JUERGEN - indent this sectoin maybe. --> ===== Baselineboard Stacking ===== In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baselineboards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline, however we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> <!-- JUERGEN - indent this sectoin maybe. --> ===== Recirculation ===== "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. ==== 3.3.1 Narrow Subbands with the 3-bit sampler ==== As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs with dual polarization, 2MHz resolution and 128MHz subbands. <!-- JUERGEN -anything else to say? Is this only dual polarization? --> === 3.4 Data Rate Limits === Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. == 4 Tips for Planning, Setup, and Processing of VLA Spectral Line Observations == '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing <!-- JUERGEN -indent this --> two 1GHz baseband pairs <!-- JUERGEN -indent this --> up to 16 independently tunable subband pairs per baseband pair <!-- JUERGEN -indent this --> independent subband bandwidths ranging from 31.25kHz to 128MHz <!-- JUERGEN -should this be 31.25kHz-128MHz? --> <!-- JUERGEN -indent this --> independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN -once again, I did 2048 channels in dual pol. Is this possible? --> * 3-bit samplers providing <!-- JUERGEN -indent this --> four 2GHz basebands with dual polarization, 128MHz subband bandwidths, and 2MHz resolution <!--JUERGEN -is this right? --> === 4.1 Considerations for Planning Subband Bandwidths and Resolution ==== 4.1.1 Bandwidths required for UV continuum subtraction ==== When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. ==== 4.1.2 Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. === 4.2 Sensitivity/Exposure Time Calculation === [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. === 4.3 The Proposal Submission Tool (PST) === The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === 4.4 Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. ==== 4.4.1 Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> ==== 4.4.2 Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== 4.4.3 Correlator Resources Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 3de9ac02bb5f68a7dae7e8659b2efa562b8025e3 1303 1302 2012-07-02T04:38:08Z Jott 8 /* 2.2 Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == 1 Introducing VLA Spectral Line Observing == The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN - I did 2048 channels in dual pol for my RSRO. Is this possible for regular observing for Aug 2012 deadline? --> * 3-bit samplers providing ** four 2GHz basebands with dual polarization, 128MHz subbands, and 2MHz resolution <!--JUERGEN -is this right? --> == 2 Line Frequencies and Velocity Conventions == === 2.1 Line Rest Frequencies === Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == 2.2 Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == 2.3 Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Therefore in order to determine the observing frequency, one must specify the frame in which the source velocity is measured. Various rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] <!-- JUERGEN - why the extra ]'s?--> == 2.4 Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accomodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == 3 The WIDAR Correlator == === 3.1 Basebands === Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independantly tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] ==== 3.1.1 Baseband Tuning Restrictions ==== The following restrictions apply to baseband tuning: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * With 3-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD <!--JUERGEN -which BD? B1D1 or B2D2? --> * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation <!-- JUERGEN - 8-bit samplers should have AC0, not A1C1, B1D1, B2D2, etc. --> === 3.2 Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands, which actually measure the spectrum, can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. The preferred method is to avoid them in your spectral setup. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === 3.3 Correlator Resources and Subband Placement === Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 BlB pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. ==== 3.3.1 Narrow Subbands with the 8-bit sampler ==== Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This means that ALL subbands, All narrow subbands can overlap each other, and every subband is independent, so that <i> different subbands may have different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> <!-- JUERGEN - indent this sectoin maybe. --> ===== Standard Subbands ===== Standard subbands allocate a single baselineboard pair to a each subband (in single polarizations) or subband pair (in dual or full polarization). Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 <!--JUERGEN 64 total subband pairs for shared risk yes? -->) for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> <!-- JUERGEN - indent this sectoin maybe. --> ===== Baselineboard Stacking ===== In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baselineboards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline, however we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> <!-- JUERGEN - indent this sectoin maybe. --> ===== Recirculation ===== "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. ==== 3.3.1 Narrow Subbands with the 3-bit sampler ==== As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs with dual polarization, 2MHz resolution and 128MHz subbands. <!-- JUERGEN -anything else to say? Is this only dual polarization? --> === 3.4 Data Rate Limits === Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. == 4 Tips for Planning, Setup, and Processing of VLA Spectral Line Observations == '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing <!-- JUERGEN -indent this --> two 1GHz baseband pairs <!-- JUERGEN -indent this --> up to 16 independently tunable subband pairs per baseband pair <!-- JUERGEN -indent this --> independent subband bandwidths ranging from 31.25kHz to 128MHz <!-- JUERGEN -should this be 31.25kHz-128MHz? --> <!-- JUERGEN -indent this --> independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN -once again, I did 2048 channels in dual pol. Is this possible? --> * 3-bit samplers providing <!-- JUERGEN -indent this --> four 2GHz basebands with dual polarization, 128MHz subband bandwidths, and 2MHz resolution <!--JUERGEN -is this right? --> === 4.1 Considerations for Planning Subband Bandwidths and Resolution ==== 4.1.1 Bandwidths required for UV continuum subtraction ==== When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. ==== 4.1.2 Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. === 4.2 Sensitivity/Exposure Time Calculation === [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. === 4.3 The Proposal Submission Tool (PST) === The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === 4.4 Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. ==== 4.4.1 Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> ==== 4.4.2 Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== 4.4.3 Correlator Resources Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> f63e49bda30a914602a8cd390f632c7253579dff 1304 1303 2012-07-02T04:43:28Z Jott 8 /* 2.3 Velocity Frames */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == 1 Introducing VLA Spectral Line Observing == The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN - I did 2048 channels in dual pol for my RSRO. Is this possible for regular observing for Aug 2012 deadline? --> * 3-bit samplers providing ** four 2GHz basebands with dual polarization, 128MHz subbands, and 2MHz resolution <!--JUERGEN -is this right? --> == 2 Line Frequencies and Velocity Conventions == === 2.1 Line Rest Frequencies === Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == 2.2 Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == 2.3 Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == 2.4 Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accomodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == 3 The WIDAR Correlator == === 3.1 Basebands === Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independantly tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] ==== 3.1.1 Baseband Tuning Restrictions ==== The following restrictions apply to baseband tuning: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * With 3-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD <!--JUERGEN -which BD? B1D1 or B2D2? --> * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation <!-- JUERGEN - 8-bit samplers should have AC0, not A1C1, B1D1, B2D2, etc. --> === 3.2 Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands, which actually measure the spectrum, can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. The preferred method is to avoid them in your spectral setup. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === 3.3 Correlator Resources and Subband Placement === Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 BlB pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. ==== 3.3.1 Narrow Subbands with the 8-bit sampler ==== Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This means that ALL subbands, All narrow subbands can overlap each other, and every subband is independent, so that <i> different subbands may have different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> <!-- JUERGEN - indent this sectoin maybe. --> ===== Standard Subbands ===== Standard subbands allocate a single baselineboard pair to a each subband (in single polarizations) or subband pair (in dual or full polarization). Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 <!--JUERGEN 64 total subband pairs for shared risk yes? -->) for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> <!-- JUERGEN - indent this sectoin maybe. --> ===== Baselineboard Stacking ===== In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baselineboards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline, however we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> <!-- JUERGEN - indent this sectoin maybe. --> ===== Recirculation ===== "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. ==== 3.3.1 Narrow Subbands with the 3-bit sampler ==== As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs with dual polarization, 2MHz resolution and 128MHz subbands. <!-- JUERGEN -anything else to say? Is this only dual polarization? --> === 3.4 Data Rate Limits === Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. == 4 Tips for Planning, Setup, and Processing of VLA Spectral Line Observations == '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing <!-- JUERGEN -indent this --> two 1GHz baseband pairs <!-- JUERGEN -indent this --> up to 16 independently tunable subband pairs per baseband pair <!-- JUERGEN -indent this --> independent subband bandwidths ranging from 31.25kHz to 128MHz <!-- JUERGEN -should this be 31.25kHz-128MHz? --> <!-- JUERGEN -indent this --> independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN -once again, I did 2048 channels in dual pol. Is this possible? --> * 3-bit samplers providing <!-- JUERGEN -indent this --> four 2GHz basebands with dual polarization, 128MHz subband bandwidths, and 2MHz resolution <!--JUERGEN -is this right? --> === 4.1 Considerations for Planning Subband Bandwidths and Resolution ==== 4.1.1 Bandwidths required for UV continuum subtraction ==== When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. ==== 4.1.2 Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. === 4.2 Sensitivity/Exposure Time Calculation === [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. === 4.3 The Proposal Submission Tool (PST) === The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === 4.4 Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. ==== 4.4.1 Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> ==== 4.4.2 Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== 4.4.3 Correlator Resources Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> b288d96066e56adc3f18128f721f99edebf129d5 1305 1304 2012-07-02T04:44:04Z Jott 8 /* 2.3 Velocity Frames */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == 1 Introducing VLA Spectral Line Observing == The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN - I did 2048 channels in dual pol for my RSRO. Is this possible for regular observing for Aug 2012 deadline? --> * 3-bit samplers providing ** four 2GHz basebands with dual polarization, 128MHz subbands, and 2MHz resolution <!--JUERGEN -is this right? --> == 2 Line Frequencies and Velocity Conventions == === 2.1 Line Rest Frequencies === Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == 2.2 Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == 2.3 Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == 2.4 Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accomodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == 3 The WIDAR Correlator == === 3.1 Basebands === Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independantly tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] ==== 3.1.1 Baseband Tuning Restrictions ==== The following restrictions apply to baseband tuning: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * With 3-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD <!--JUERGEN -which BD? B1D1 or B2D2? --> * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation <!-- JUERGEN - 8-bit samplers should have AC0, not A1C1, B1D1, B2D2, etc. --> === 3.2 Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands, which actually measure the spectrum, can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. The preferred method is to avoid them in your spectral setup. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === 3.3 Correlator Resources and Subband Placement === Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 BlB pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. ==== 3.3.1 Narrow Subbands with the 8-bit sampler ==== Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This means that ALL subbands, All narrow subbands can overlap each other, and every subband is independent, so that <i> different subbands may have different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> <!-- JUERGEN - indent this sectoin maybe. --> ===== Standard Subbands ===== Standard subbands allocate a single baselineboard pair to a each subband (in single polarizations) or subband pair (in dual or full polarization). Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 <!--JUERGEN 64 total subband pairs for shared risk yes? -->) for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> <!-- JUERGEN - indent this sectoin maybe. --> ===== Baselineboard Stacking ===== In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baselineboards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline, however we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> <!-- JUERGEN - indent this sectoin maybe. --> ===== Recirculation ===== "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. ==== 3.3.1 Narrow Subbands with the 3-bit sampler ==== As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs with dual polarization, 2MHz resolution and 128MHz subbands. <!-- JUERGEN -anything else to say? Is this only dual polarization? --> === 3.4 Data Rate Limits === Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. == 4 Tips for Planning, Setup, and Processing of VLA Spectral Line Observations == '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing <!-- JUERGEN -indent this --> two 1GHz baseband pairs <!-- JUERGEN -indent this --> up to 16 independently tunable subband pairs per baseband pair <!-- JUERGEN -indent this --> independent subband bandwidths ranging from 31.25kHz to 128MHz <!-- JUERGEN -should this be 31.25kHz-128MHz? --> <!-- JUERGEN -indent this --> independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN -once again, I did 2048 channels in dual pol. Is this possible? --> * 3-bit samplers providing <!-- JUERGEN -indent this --> four 2GHz basebands with dual polarization, 128MHz subband bandwidths, and 2MHz resolution <!--JUERGEN -is this right? --> === 4.1 Considerations for Planning Subband Bandwidths and Resolution ==== 4.1.1 Bandwidths required for UV continuum subtraction ==== When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. ==== 4.1.2 Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. === 4.2 Sensitivity/Exposure Time Calculation === [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. === 4.3 The Proposal Submission Tool (PST) === The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === 4.4 Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. ==== 4.4.1 Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> ==== 4.4.2 Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== 4.4.3 Correlator Resources Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 14a2940eff8ed4e899490bfb74be6afe0dd6e7c2 1306 1305 2012-07-02T04:45:36Z Jott 8 /* 2.4 Doppler Correction */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == 1 Introducing VLA Spectral Line Observing == The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN - I did 2048 channels in dual pol for my RSRO. Is this possible for regular observing for Aug 2012 deadline? --> * 3-bit samplers providing ** four 2GHz basebands with dual polarization, 128MHz subbands, and 2MHz resolution <!--JUERGEN -is this right? --> == 2 Line Frequencies and Velocity Conventions == === 2.1 Line Rest Frequencies === Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == 2.2 Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == 2.3 Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == 2.4 Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == 3 The WIDAR Correlator == === 3.1 Basebands === Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independantly tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] ==== 3.1.1 Baseband Tuning Restrictions ==== The following restrictions apply to baseband tuning: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * With 3-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD <!--JUERGEN -which BD? B1D1 or B2D2? --> * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation <!-- JUERGEN - 8-bit samplers should have AC0, not A1C1, B1D1, B2D2, etc. --> === 3.2 Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands, which actually measure the spectrum, can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. The preferred method is to avoid them in your spectral setup. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === 3.3 Correlator Resources and Subband Placement === Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 BlB pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. ==== 3.3.1 Narrow Subbands with the 8-bit sampler ==== Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This means that ALL subbands, All narrow subbands can overlap each other, and every subband is independent, so that <i> different subbands may have different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> <!-- JUERGEN - indent this sectoin maybe. --> ===== Standard Subbands ===== Standard subbands allocate a single baselineboard pair to a each subband (in single polarizations) or subband pair (in dual or full polarization). Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 <!--JUERGEN 64 total subband pairs for shared risk yes? -->) for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> <!-- JUERGEN - indent this sectoin maybe. --> ===== Baselineboard Stacking ===== In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baselineboards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline, however we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> <!-- JUERGEN - indent this sectoin maybe. --> ===== Recirculation ===== "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. ==== 3.3.1 Narrow Subbands with the 3-bit sampler ==== As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs with dual polarization, 2MHz resolution and 128MHz subbands. <!-- JUERGEN -anything else to say? Is this only dual polarization? --> === 3.4 Data Rate Limits === Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. == 4 Tips for Planning, Setup, and Processing of VLA Spectral Line Observations == '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing <!-- JUERGEN -indent this --> two 1GHz baseband pairs <!-- JUERGEN -indent this --> up to 16 independently tunable subband pairs per baseband pair <!-- JUERGEN -indent this --> independent subband bandwidths ranging from 31.25kHz to 128MHz <!-- JUERGEN -should this be 31.25kHz-128MHz? --> <!-- JUERGEN -indent this --> independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN -once again, I did 2048 channels in dual pol. Is this possible? --> * 3-bit samplers providing <!-- JUERGEN -indent this --> four 2GHz basebands with dual polarization, 128MHz subband bandwidths, and 2MHz resolution <!--JUERGEN -is this right? --> === 4.1 Considerations for Planning Subband Bandwidths and Resolution ==== 4.1.1 Bandwidths required for UV continuum subtraction ==== When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. ==== 4.1.2 Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. === 4.2 Sensitivity/Exposure Time Calculation === [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. === 4.3 The Proposal Submission Tool (PST) === The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === 4.4 Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. ==== 4.4.1 Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> ==== 4.4.2 Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== 4.4.3 Correlator Resources Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> b9aaff4c1c0cd1eb05412c7382348f46dcd67b1f 1307 1306 2012-07-02T04:48:37Z Jott 8 /* 3.1 Basebands */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == 1 Introducing VLA Spectral Line Observing == The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN - I did 2048 channels in dual pol for my RSRO. Is this possible for regular observing for Aug 2012 deadline? --> * 3-bit samplers providing ** four 2GHz basebands with dual polarization, 128MHz subbands, and 2MHz resolution <!--JUERGEN -is this right? --> == 2 Line Frequencies and Velocity Conventions == === 2.1 Line Rest Frequencies === Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == 2.2 Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == 2.3 Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == 2.4 Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == 3 The WIDAR Correlator == === 3.1 Basebands === Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] ==== 3.1.1 Baseband Tuning Restrictions ==== The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation === 3.2 Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands, which actually measure the spectrum, can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. The preferred method is to avoid them in your spectral setup. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === 3.3 Correlator Resources and Subband Placement === Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 BlB pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. ==== 3.3.1 Narrow Subbands with the 8-bit sampler ==== Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This means that ALL subbands, All narrow subbands can overlap each other, and every subband is independent, so that <i> different subbands may have different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> <!-- JUERGEN - indent this sectoin maybe. --> ===== Standard Subbands ===== Standard subbands allocate a single baselineboard pair to a each subband (in single polarizations) or subband pair (in dual or full polarization). Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 <!--JUERGEN 64 total subband pairs for shared risk yes? -->) for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> <!-- JUERGEN - indent this sectoin maybe. --> ===== Baselineboard Stacking ===== In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baselineboards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline, however we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> <!-- JUERGEN - indent this sectoin maybe. --> ===== Recirculation ===== "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. ==== 3.3.1 Narrow Subbands with the 3-bit sampler ==== As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs with dual polarization, 2MHz resolution and 128MHz subbands. <!-- JUERGEN -anything else to say? Is this only dual polarization? --> === 3.4 Data Rate Limits === Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. == 4 Tips for Planning, Setup, and Processing of VLA Spectral Line Observations == '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing <!-- JUERGEN -indent this --> two 1GHz baseband pairs <!-- JUERGEN -indent this --> up to 16 independently tunable subband pairs per baseband pair <!-- JUERGEN -indent this --> independent subband bandwidths ranging from 31.25kHz to 128MHz <!-- JUERGEN -should this be 31.25kHz-128MHz? --> <!-- JUERGEN -indent this --> independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN -once again, I did 2048 channels in dual pol. Is this possible? --> * 3-bit samplers providing <!-- JUERGEN -indent this --> four 2GHz basebands with dual polarization, 128MHz subband bandwidths, and 2MHz resolution <!--JUERGEN -is this right? --> === 4.1 Considerations for Planning Subband Bandwidths and Resolution ==== 4.1.1 Bandwidths required for UV continuum subtraction ==== When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. ==== 4.1.2 Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. === 4.2 Sensitivity/Exposure Time Calculation === [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. === 4.3 The Proposal Submission Tool (PST) === The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === 4.4 Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. ==== 4.4.1 Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> ==== 4.4.2 Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== 4.4.3 Correlator Resources Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 0d577abeb732a852ce2a2918ddd1f31a9d2850c7 1308 1307 2012-07-02T05:02:49Z Jott 8 /* 3.2 Fixed 128MHz Subbands and 128MHz "Suckouts" */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == 1 Introducing VLA Spectral Line Observing == The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN - I did 2048 channels in dual pol for my RSRO. Is this possible for regular observing for Aug 2012 deadline? --> * 3-bit samplers providing ** four 2GHz basebands with dual polarization, 128MHz subbands, and 2MHz resolution <!--JUERGEN -is this right? --> == 2 Line Frequencies and Velocity Conventions == === 2.1 Line Rest Frequencies === Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == 2.2 Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == 2.3 Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == 2.4 Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == 3 The WIDAR Correlator == === 3.1 Basebands === Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] ==== 3.1.1 Baseband Tuning Restrictions ==== The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation === 3.2 Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === 3.3 Correlator Resources and Subband Placement === Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 BlB pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. ==== 3.3.1 Narrow Subbands with the 8-bit sampler ==== Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This means that ALL subbands, All narrow subbands can overlap each other, and every subband is independent, so that <i> different subbands may have different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> <!-- JUERGEN - indent this sectoin maybe. --> ===== Standard Subbands ===== Standard subbands allocate a single baselineboard pair to a each subband (in single polarizations) or subband pair (in dual or full polarization). Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 <!--JUERGEN 64 total subband pairs for shared risk yes? -->) for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> <!-- JUERGEN - indent this sectoin maybe. --> ===== Baselineboard Stacking ===== In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baselineboards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline, however we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> <!-- JUERGEN - indent this sectoin maybe. --> ===== Recirculation ===== "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. ==== 3.3.1 Narrow Subbands with the 3-bit sampler ==== As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs with dual polarization, 2MHz resolution and 128MHz subbands. <!-- JUERGEN -anything else to say? Is this only dual polarization? --> === 3.4 Data Rate Limits === Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. == 4 Tips for Planning, Setup, and Processing of VLA Spectral Line Observations == '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing <!-- JUERGEN -indent this --> two 1GHz baseband pairs <!-- JUERGEN -indent this --> up to 16 independently tunable subband pairs per baseband pair <!-- JUERGEN -indent this --> independent subband bandwidths ranging from 31.25kHz to 128MHz <!-- JUERGEN -should this be 31.25kHz-128MHz? --> <!-- JUERGEN -indent this --> independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN -once again, I did 2048 channels in dual pol. Is this possible? --> * 3-bit samplers providing <!-- JUERGEN -indent this --> four 2GHz basebands with dual polarization, 128MHz subband bandwidths, and 2MHz resolution <!--JUERGEN -is this right? --> === 4.1 Considerations for Planning Subband Bandwidths and Resolution ==== 4.1.1 Bandwidths required for UV continuum subtraction ==== When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. ==== 4.1.2 Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. === 4.2 Sensitivity/Exposure Time Calculation === [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. === 4.3 The Proposal Submission Tool (PST) === The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === 4.4 Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. ==== 4.4.1 Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> ==== 4.4.2 Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== 4.4.3 Correlator Resources Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 3d618ef4dfb4cc96434d101666c92bc42ae52375 1309 1308 2012-07-02T05:18:20Z Jott 8 /* 3.3 Correlator Resources and Subband Placement */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == 1 Introducing VLA Spectral Line Observing == The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN - I did 2048 channels in dual pol for my RSRO. Is this possible for regular observing for Aug 2012 deadline? --> * 3-bit samplers providing ** four 2GHz basebands with dual polarization, 128MHz subbands, and 2MHz resolution <!--JUERGEN -is this right? --> == 2 Line Frequencies and Velocity Conventions == === 2.1 Line Rest Frequencies === Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == 2.2 Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == 2.3 Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == 2.4 Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == 3 The WIDAR Correlator == === 3.1 Basebands === Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] ==== 3.1.1 Baseband Tuning Restrictions ==== The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation === 3.2 Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === 3.3 Correlator Resources and Subband Placement === Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. ==== 3.3.1 Narrow Subbands with the 8-bit sampler ==== Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> ===== Standard Subbands ===== Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ===== Baselineboard Stacking ===== In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> ===== Recirculation ===== "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. ==== 3.3.1 Narrow Subbands with the 3-bit sampler ==== As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs with full polarization, 2MHz resolution and 128MHz subbands. === 3.4 Data Rate Limits === Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. == 4 Tips for Planning, Setup, and Processing of VLA Spectral Line Observations == '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing <!-- JUERGEN -indent this --> two 1GHz baseband pairs <!-- JUERGEN -indent this --> up to 16 independently tunable subband pairs per baseband pair <!-- JUERGEN -indent this --> independent subband bandwidths ranging from 31.25kHz to 128MHz <!-- JUERGEN -should this be 31.25kHz-128MHz? --> <!-- JUERGEN -indent this --> independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN -once again, I did 2048 channels in dual pol. Is this possible? --> * 3-bit samplers providing <!-- JUERGEN -indent this --> four 2GHz basebands with dual polarization, 128MHz subband bandwidths, and 2MHz resolution <!--JUERGEN -is this right? --> === 4.1 Considerations for Planning Subband Bandwidths and Resolution ==== 4.1.1 Bandwidths required for UV continuum subtraction ==== When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. ==== 4.1.2 Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. === 4.2 Sensitivity/Exposure Time Calculation === [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. === 4.3 The Proposal Submission Tool (PST) === The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === 4.4 Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. ==== 4.4.1 Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> ==== 4.4.2 Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== 4.4.3 Correlator Resources Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 0c8b0857fa7c14d5b7094870d14566702b02b5b0 1310 1309 2012-07-02T05:21:17Z Jott 8 /* 3.3.1 Narrow Subbands with the 3-bit sampler */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == 1 Introducing VLA Spectral Line Observing == The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN - I did 2048 channels in dual pol for my RSRO. Is this possible for regular observing for Aug 2012 deadline? --> * 3-bit samplers providing ** four 2GHz basebands with dual polarization, 128MHz subbands, and 2MHz resolution <!--JUERGEN -is this right? --> == 2 Line Frequencies and Velocity Conventions == === 2.1 Line Rest Frequencies === Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == 2.2 Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == 2.3 Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == 2.4 Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == 3 The WIDAR Correlator == === 3.1 Basebands === Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] ==== 3.1.1 Baseband Tuning Restrictions ==== The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation === 3.2 Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === 3.3 Correlator Resources and Subband Placement === Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. ==== 3.3.1 Narrow Subbands with the 8-bit sampler ==== Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> ===== Standard Subbands ===== Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ===== Baselineboard Stacking ===== In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> ===== Recirculation ===== "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. ==== 3.3.1 Narrow Subbands with the 3-bit sampler ==== As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. === 3.4 Data Rate Limits === Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. == 4 Tips for Planning, Setup, and Processing of VLA Spectral Line Observations == '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing <!-- JUERGEN -indent this --> two 1GHz baseband pairs <!-- JUERGEN -indent this --> up to 16 independently tunable subband pairs per baseband pair <!-- JUERGEN -indent this --> independent subband bandwidths ranging from 31.25kHz to 128MHz <!-- JUERGEN -should this be 31.25kHz-128MHz? --> <!-- JUERGEN -indent this --> independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN -once again, I did 2048 channels in dual pol. Is this possible? --> * 3-bit samplers providing <!-- JUERGEN -indent this --> four 2GHz basebands with dual polarization, 128MHz subband bandwidths, and 2MHz resolution <!--JUERGEN -is this right? --> === 4.1 Considerations for Planning Subband Bandwidths and Resolution ==== 4.1.1 Bandwidths required for UV continuum subtraction ==== When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. ==== 4.1.2 Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. === 4.2 Sensitivity/Exposure Time Calculation === [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. === 4.3 The Proposal Submission Tool (PST) === The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === 4.4 Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. ==== 4.4.1 Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> ==== 4.4.2 Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== 4.4.3 Correlator Resources Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> c7fbce65b0fbc631981e1fdb820e3fb347821051 1311 1310 2012-07-02T05:22:23Z Jott 8 /* Standard Subbands */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == 1 Introducing VLA Spectral Line Observing == The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN - I did 2048 channels in dual pol for my RSRO. Is this possible for regular observing for Aug 2012 deadline? --> * 3-bit samplers providing ** four 2GHz basebands with dual polarization, 128MHz subbands, and 2MHz resolution <!--JUERGEN -is this right? --> == 2 Line Frequencies and Velocity Conventions == === 2.1 Line Rest Frequencies === Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == 2.2 Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == 2.3 Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == 2.4 Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == 3 The WIDAR Correlator == === 3.1 Basebands === Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] ==== 3.1.1 Baseband Tuning Restrictions ==== The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation === 3.2 Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === 3.3 Correlator Resources and Subband Placement === Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. ==== 3.3.1 Narrow Subbands with the 8-bit sampler ==== Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> ==== Standard Subbands ==== Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ===== Baselineboard Stacking ===== In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> ===== Recirculation ===== "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. ==== 3.3.1 Narrow Subbands with the 3-bit sampler ==== As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. === 3.4 Data Rate Limits === Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. == 4 Tips for Planning, Setup, and Processing of VLA Spectral Line Observations == '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing <!-- JUERGEN -indent this --> two 1GHz baseband pairs <!-- JUERGEN -indent this --> up to 16 independently tunable subband pairs per baseband pair <!-- JUERGEN -indent this --> independent subband bandwidths ranging from 31.25kHz to 128MHz <!-- JUERGEN -should this be 31.25kHz-128MHz? --> <!-- JUERGEN -indent this --> independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN -once again, I did 2048 channels in dual pol. Is this possible? --> * 3-bit samplers providing <!-- JUERGEN -indent this --> four 2GHz basebands with dual polarization, 128MHz subband bandwidths, and 2MHz resolution <!--JUERGEN -is this right? --> === 4.1 Considerations for Planning Subband Bandwidths and Resolution ==== 4.1.1 Bandwidths required for UV continuum subtraction ==== When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. ==== 4.1.2 Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. === 4.2 Sensitivity/Exposure Time Calculation === [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. === 4.3 The Proposal Submission Tool (PST) === The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === 4.4 Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. ==== 4.4.1 Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> ==== 4.4.2 Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== 4.4.3 Correlator Resources Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 923541d17c593724941302c1abbcbd497ef29774 1312 1311 2012-07-02T05:23:51Z Jott 8 /* Standard Subbands */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == 1 Introducing VLA Spectral Line Observing == The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN - I did 2048 channels in dual pol for my RSRO. Is this possible for regular observing for Aug 2012 deadline? --> * 3-bit samplers providing ** four 2GHz basebands with dual polarization, 128MHz subbands, and 2MHz resolution <!--JUERGEN -is this right? --> == 2 Line Frequencies and Velocity Conventions == === 2.1 Line Rest Frequencies === Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == 2.2 Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == 2.3 Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == 2.4 Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == 3 The WIDAR Correlator == === 3.1 Basebands === Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] ==== 3.1.1 Baseband Tuning Restrictions ==== The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation === 3.2 Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === 3.3 Correlator Resources and Subband Placement === Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. ==== 3.3.1 Narrow Subbands with the 8-bit sampler ==== Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> ==== Standard Subbands ==== Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== Baselineboard Stacking ==== In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> ==== Recirculation ==== "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. ==== 3.3.1 Narrow Subbands with the 3-bit sampler ==== As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. === 3.4 Data Rate Limits === Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. == 4 Tips for Planning, Setup, and Processing of VLA Spectral Line Observations == '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing <!-- JUERGEN -indent this --> two 1GHz baseband pairs <!-- JUERGEN -indent this --> up to 16 independently tunable subband pairs per baseband pair <!-- JUERGEN -indent this --> independent subband bandwidths ranging from 31.25kHz to 128MHz <!-- JUERGEN -should this be 31.25kHz-128MHz? --> <!-- JUERGEN -indent this --> independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN -once again, I did 2048 channels in dual pol. Is this possible? --> * 3-bit samplers providing <!-- JUERGEN -indent this --> four 2GHz basebands with dual polarization, 128MHz subband bandwidths, and 2MHz resolution <!--JUERGEN -is this right? --> === 4.1 Considerations for Planning Subband Bandwidths and Resolution ==== 4.1.1 Bandwidths required for UV continuum subtraction ==== When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. ==== 4.1.2 Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. === 4.2 Sensitivity/Exposure Time Calculation === [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. === 4.3 The Proposal Submission Tool (PST) === The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === 4.4 Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. ==== 4.4.1 Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> ==== 4.4.2 Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== 4.4.3 Correlator Resources Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 0486626d45b030dceaf245ad7d44973cbc9be241 1313 1312 2012-07-02T05:24:24Z Jott 8 /* 3.3.1 Narrow Subbands with the 8-bit sampler */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == 1 Introducing VLA Spectral Line Observing == The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN - I did 2048 channels in dual pol for my RSRO. Is this possible for regular observing for Aug 2012 deadline? --> * 3-bit samplers providing ** four 2GHz basebands with dual polarization, 128MHz subbands, and 2MHz resolution <!--JUERGEN -is this right? --> == 2 Line Frequencies and Velocity Conventions == === 2.1 Line Rest Frequencies === Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == 2.2 Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == 2.3 Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == 2.4 Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == 3 The WIDAR Correlator == === 3.1 Basebands === Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] ==== 3.1.1 Baseband Tuning Restrictions ==== The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation === 3.2 Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === 3.3 Correlator Resources and Subband Placement === Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. ==== 3.3.1 Narrow Subbands with the 8-bit sampler ==== Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> ==== Standard Subbands ==== Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== Baselineboard Stacking ==== In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> ==== Recirculation ==== "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. ==== 3.3.1 Narrow Subbands with the 3-bit sampler ==== As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. === 3.4 Data Rate Limits === Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. == 4 Tips for Planning, Setup, and Processing of VLA Spectral Line Observations == '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing <!-- JUERGEN -indent this --> two 1GHz baseband pairs <!-- JUERGEN -indent this --> up to 16 independently tunable subband pairs per baseband pair <!-- JUERGEN -indent this --> independent subband bandwidths ranging from 31.25kHz to 128MHz <!-- JUERGEN -should this be 31.25kHz-128MHz? --> <!-- JUERGEN -indent this --> independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN -once again, I did 2048 channels in dual pol. Is this possible? --> * 3-bit samplers providing <!-- JUERGEN -indent this --> four 2GHz basebands with dual polarization, 128MHz subband bandwidths, and 2MHz resolution <!--JUERGEN -is this right? --> === 4.1 Considerations for Planning Subband Bandwidths and Resolution ==== 4.1.1 Bandwidths required for UV continuum subtraction ==== When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. ==== 4.1.2 Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. === 4.2 Sensitivity/Exposure Time Calculation === [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. === 4.3 The Proposal Submission Tool (PST) === The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === 4.4 Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. ==== 4.4.1 Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> ==== 4.4.2 Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== 4.4.3 Correlator Resources Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 362aafefcc3c1c174f86da1429ce319063bdce61 1314 1313 2012-07-02T05:29:59Z Jott 8 wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> = Introducing VLA Spectral Line Observing = The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN - I did 2048 channels in dual pol for my RSRO. Is this possible for regular observing for Aug 2012 deadline? --> * 3-bit samplers providing ** four 2GHz basebands with dual polarization, 128MHz subbands, and 2MHz resolution <!--JUERGEN -is this right? --> = Line Frequencies and Velocity Conventions = == Line Rest Frequencies == Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. = The WIDAR Correlator = == Basebands == Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Baseband Tuning Restrictions === The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation == Fixed 128MHz Subbands and 128MHz "Suckouts" == After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] == Correlator Resources and Subband Placement == Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. === Narrow Subbands with the 8-bit sampler === Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> === Standard Subbands === Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> === Recirculation === "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. == Narrow Subbands with the 3-bit sampler == As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. == Data Rate Limits == Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. = Tips for Planning, Setup, and Processing of VLA Spectral Line Observations = '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing <!-- JUERGEN -indent this --> two 1GHz baseband pairs <!-- JUERGEN -indent this --> up to 16 independently tunable subband pairs per baseband pair <!-- JUERGEN -indent this --> independent subband bandwidths ranging from 31.25kHz to 128MHz <!-- JUERGEN -should this be 31.25kHz-128MHz? --> <!-- JUERGEN -indent this --> independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN -once again, I did 2048 channels in dual pol. Is this possible? --> * 3-bit samplers providing <!-- JUERGEN -indent this --> four 2GHz basebands with dual polarization, 128MHz subband bandwidths, and 2MHz resolution <!--JUERGEN -is this right? --> == Considerations for Planning Subband Bandwidths and Resolution == === Bandwidths required for UV continuum subtraction === When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. === Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The Proposal Submission Tool (PST) == The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. == Setting up a Spectral Observation using the Observation Preparation Tool (OPT) == The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. === Bandpass Setup === All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> === Phase/Complex Gain Calibration === The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. === Correlator Resources Setup === '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 69f8554e0e14941b084a5730e688cac77bd6715a 1315 1314 2012-07-02T05:33:25Z Jott 8 /* Tips for Planning, Setup, and Processing of VLA Spectral Line Observations */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> = Introducing VLA Spectral Line Observing = The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels <!-- JUERGEN - I did 2048 channels in dual pol for my RSRO. Is this possible for regular observing for Aug 2012 deadline? --> * 3-bit samplers providing ** four 2GHz basebands with dual polarization, 128MHz subbands, and 2MHz resolution <!--JUERGEN -is this right? --> = Line Frequencies and Velocity Conventions = == Line Rest Frequencies == Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. = The WIDAR Correlator = == Basebands == Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Baseband Tuning Restrictions === The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation == Fixed 128MHz Subbands and 128MHz "Suckouts" == After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] == Correlator Resources and Subband Placement == Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. === Narrow Subbands with the 8-bit sampler === Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> === Standard Subbands === Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> === Recirculation === "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. == Narrow Subbands with the 3-bit sampler == As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. == Data Rate Limits == Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. = Tips for Planning, Setup, and Processing of VLA Spectral Line Observations = '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz baseband pairs ** up to 16 independently tunable subband pairs per baseband pair ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization == Considerations for Planning Subband Bandwidths and Resolution == === Bandwidths required for UV continuum subtraction === When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. === Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The Proposal Submission Tool (PST) == The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. == Setting up a Spectral Observation using the Observation Preparation Tool (OPT) == The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. === Bandpass Setup === All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> === Phase/Complex Gain Calibration === The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. === Correlator Resources Setup === '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 67a59d638a9dd104f8cb48cd4f9ffdeb7502906a 1316 1315 2012-07-02T05:36:54Z Jott 8 /* Introducing VLA Spectral Line Observing */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> = Introducing VLA Spectral Line Observing = The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization = Line Frequencies and Velocity Conventions = == Line Rest Frequencies == Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. = The WIDAR Correlator = == Basebands == Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Baseband Tuning Restrictions === The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation == Fixed 128MHz Subbands and 128MHz "Suckouts" == After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] == Correlator Resources and Subband Placement == Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. === Narrow Subbands with the 8-bit sampler === Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> === Standard Subbands === Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> === Recirculation === "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. == Narrow Subbands with the 3-bit sampler == As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. == Data Rate Limits == Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. = Tips for Planning, Setup, and Processing of VLA Spectral Line Observations = '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 15MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz baseband pairs ** up to 16 independently tunable subband pairs per baseband pair ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization == Considerations for Planning Subband Bandwidths and Resolution == === Bandwidths required for UV continuum subtraction === When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. === Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The Proposal Submission Tool (PST) == The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. == Setting up a Spectral Observation using the Observation Preparation Tool (OPT) == The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. === Bandpass Setup === All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> === Phase/Complex Gain Calibration === The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. === Correlator Resources Setup === '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> e0f1b2af36d633c5a86e6e58bd10f0981fb09e6a 1317 1316 2012-07-02T05:37:47Z Jott 8 /* Tips for Planning, Setup, and Processing of VLA Spectral Line Observations */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> = Introducing VLA Spectral Line Observing = The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization = Line Frequencies and Velocity Conventions = == Line Rest Frequencies == Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. = The WIDAR Correlator = == Basebands == Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Baseband Tuning Restrictions === The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation == Fixed 128MHz Subbands and 128MHz "Suckouts" == After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] == Correlator Resources and Subband Placement == Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. === Narrow Subbands with the 8-bit sampler === Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> === Standard Subbands === Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> === Recirculation === "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. == Narrow Subbands with the 3-bit sampler == As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. == Data Rate Limits == Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and 60MB/s <!--JUERGEN, this is correct, yes? --> for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. = Tips for Planning, Setup, and Processing of VLA Spectral Line Observations = '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization == Considerations for Planning Subband Bandwidths and Resolution == === Bandwidths required for UV continuum subtraction === When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. === Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The Proposal Submission Tool (PST) == The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. == Setting up a Spectral Observation using the Observation Preparation Tool (OPT) == The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. === Bandpass Setup === All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> === Phase/Complex Gain Calibration === The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. === Correlator Resources Setup === '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> ab236100ebe06b6a164946c5b9022264287812bd 1318 1317 2012-07-02T05:40:41Z Jott 8 /* Data Rate Limits */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> = Introducing VLA Spectral Line Observing = The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization = Line Frequencies and Velocity Conventions = == Line Rest Frequencies == Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. = The WIDAR Correlator = == Basebands == Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Baseband Tuning Restrictions === The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation == Fixed 128MHz Subbands and 128MHz "Suckouts" == After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] == Correlator Resources and Subband Placement == Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. === Narrow Subbands with the 8-bit sampler === Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> === Standard Subbands === Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> === Recirculation === "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. == Narrow Subbands with the 3-bit sampler == As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. == Data Rate Limits == Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 20MB/s for regular and more for shared risk observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary for details]. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. = Tips for Planning, Setup, and Processing of VLA Spectral Line Observations = '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization == Considerations for Planning Subband Bandwidths and Resolution == === Bandwidths required for UV continuum subtraction === When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. === Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the old VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon prominently and Hanning smoothing was frequently applied online during the observations. The Gibbs phenomenon is much less common for the upgraded VLA due to the WIDAR correlator's better design. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. Consequently, the VLA does not support online Hanning smoothing. Gibbs ringing can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The Proposal Submission Tool (PST) == The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. == Setting up a Spectral Observation using the Observation Preparation Tool (OPT) == The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. === Bandpass Setup === All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> === Phase/Complex Gain Calibration === The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. === Correlator Resources Setup === '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> e5ddb8a0f6ce07beb30374022f5e5db04303e167 0 48 1319 1318 2012-07-02T05:53:31Z Jott 8 /* Considerations for Planning Subband Bandwidths and Resolution */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> = Introducing VLA Spectral Line Observing = The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization = Line Frequencies and Velocity Conventions = == Line Rest Frequencies == Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. = The WIDAR Correlator = == Basebands == Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Baseband Tuning Restrictions === The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation == Fixed 128MHz Subbands and 128MHz "Suckouts" == After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] == Correlator Resources and Subband Placement == Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. === Narrow Subbands with the 8-bit sampler === Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> === Standard Subbands === Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> === Recirculation === "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. == Narrow Subbands with the 3-bit sampler == As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. == Data Rate Limits == Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 20MB/s for regular and more for shared risk observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary for details]. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. = Tips for Planning, Setup, and Processing of VLA Spectral Line Observations = '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization == Considerations for Planning Subband Bandwidths and Resolution == === Bandwidths required for UV continuum subtraction === When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. === Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon, a sinc function that zig-zags on alternating channels. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing is the most effective method, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the pre-upgrade VLA days, the correlator design had a realtively short, truncated lag spectrum, which could result in prominent Gibbs ringing. To avoid this effect, Hanning smoothing was frequently applied online during the observations. With the new WIDAR capabilities of the upgraded VLA, however, ringing is very rare and only observed for extremely strong maser or RFI sources. Consequently, the VLA does not support online Hanning smoothing anymore; if required, Hanning smoothing can be applied during post-processing (e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>].). The Gibbs effect can also be reduced by using higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reducing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required <!-- JUERGEN - this is rms noise, yes? --> rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The Proposal Submission Tool (PST) == The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. == Setting up a Spectral Observation using the Observation Preparation Tool (OPT) == The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. === Bandpass Setup === All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> === Phase/Complex Gain Calibration === The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. === Correlator Resources Setup === '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> e16e11da3413dce2455cd25a0e3c07e5b9500da9 1320 1319 2012-07-02T05:55:25Z Jott 8 /* Sensitivity/Exposure Time Calculation */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> = Introducing VLA Spectral Line Observing = The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization = Line Frequencies and Velocity Conventions = == Line Rest Frequencies == Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. = The WIDAR Correlator = == Basebands == Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Baseband Tuning Restrictions === The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation == Fixed 128MHz Subbands and 128MHz "Suckouts" == After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] == Correlator Resources and Subband Placement == Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. === Narrow Subbands with the 8-bit sampler === Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> === Standard Subbands === Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> === Recirculation === "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. == Narrow Subbands with the 3-bit sampler == As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. == Data Rate Limits == Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 20MB/s for regular and more for shared risk observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary for details]. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. = Tips for Planning, Setup, and Processing of VLA Spectral Line Observations = '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization == Considerations for Planning Subband Bandwidths and Resolution == === Bandwidths required for UV continuum subtraction === When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. === Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon, a sinc function that zig-zags on alternating channels. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing is the most effective method, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the pre-upgrade VLA days, the correlator design had a realtively short, truncated lag spectrum, which could result in prominent Gibbs ringing. To avoid this effect, Hanning smoothing was frequently applied online during the observations. With the new WIDAR capabilities of the upgraded VLA, however, ringing is very rare and only observed for extremely strong maser or RFI sources. Consequently, the VLA does not support online Hanning smoothing anymore; if required, Hanning smoothing can be applied during post-processing (e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>].). The Gibbs effect can also be reduced by using higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reducing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on-source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and number of polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The Proposal Submission Tool (PST) == The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool, which is described in full at ''' --- ??? link to Michael's manual'''. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. == Setting up a Spectral Observation using the Observation Preparation Tool (OPT) == The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. === Bandpass Setup === All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> === Phase/Complex Gain Calibration === The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. === Correlator Resources Setup === '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 0fbf5c57c632f96c699227cd02de1a9ec49a7d65 1321 1320 2012-07-02T05:59:17Z Jott 8 /* The Proposal Submission Tool (PST) */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> = Introducing VLA Spectral Line Observing = The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization = Line Frequencies and Velocity Conventions = == Line Rest Frequencies == Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. = The WIDAR Correlator = == Basebands == Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Baseband Tuning Restrictions === The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation == Fixed 128MHz Subbands and 128MHz "Suckouts" == After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] == Correlator Resources and Subband Placement == Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. === Narrow Subbands with the 8-bit sampler === Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> === Standard Subbands === Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> === Recirculation === "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. == Narrow Subbands with the 3-bit sampler == As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. == Data Rate Limits == Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 20MB/s for regular and more for shared risk observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary for details]. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. = Tips for Planning, Setup, and Processing of VLA Spectral Line Observations = '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization == Considerations for Planning Subband Bandwidths and Resolution == === Bandwidths required for UV continuum subtraction === When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. === Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon, a sinc function that zig-zags on alternating channels. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing is the most effective method, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the pre-upgrade VLA days, the correlator design had a realtively short, truncated lag spectrum, which could result in prominent Gibbs ringing. To avoid this effect, Hanning smoothing was frequently applied online during the observations. With the new WIDAR capabilities of the upgraded VLA, however, ringing is very rare and only observed for extremely strong maser or RFI sources. Consequently, the VLA does not support online Hanning smoothing anymore; if required, Hanning smoothing can be applied during post-processing (e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>].). The Gibbs effect can also be reduced by using higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reducing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on-source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and number of polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The Proposal Submission Tool (PST) == The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. <!-- link to Michaels docs--> [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used for a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz for full, 1MHz for dual and 0.5MHz for single polarization products. No narrow subbands can be chosen. == Setting up a Spectral Observation using the Observation Preparation Tool (OPT) == The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the OPT QuickStart Guide <!-- JUERGEN - add the link to the quick start guide -->, and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. === Bandpass Setup === All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within the SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> === Phase/Complex Gain Calibration === The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. === Correlator Resources Setup === '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line onbservations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> <!-- JUERGEN - indent this section --> ===== Using Doppler Setting ===== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> dc27e84dcb74b37af5718e54758e03a413060673 1322 1321 2012-07-02T06:05:48Z Jott 8 /* Setting up a Spectral Observation using the Observation Preparation Tool (OPT) */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> = Introducing VLA Spectral Line Observing = The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization = Line Frequencies and Velocity Conventions = == Line Rest Frequencies == Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. = The WIDAR Correlator = == Basebands == Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Baseband Tuning Restrictions === The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation == Fixed 128MHz Subbands and 128MHz "Suckouts" == After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] == Correlator Resources and Subband Placement == Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. === Narrow Subbands with the 8-bit sampler === Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> === Standard Subbands === Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> === Recirculation === "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. == Narrow Subbands with the 3-bit sampler == As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. == Data Rate Limits == Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 20MB/s for regular and more for shared risk observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary for details]. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. = Tips for Planning, Setup, and Processing of VLA Spectral Line Observations = '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization == Considerations for Planning Subband Bandwidths and Resolution == === Bandwidths required for UV continuum subtraction === When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. === Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon, a sinc function that zig-zags on alternating channels. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing is the most effective method, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the pre-upgrade VLA days, the correlator design had a realtively short, truncated lag spectrum, which could result in prominent Gibbs ringing. To avoid this effect, Hanning smoothing was frequently applied online during the observations. With the new WIDAR capabilities of the upgraded VLA, however, ringing is very rare and only observed for extremely strong maser or RFI sources. Consequently, the VLA does not support online Hanning smoothing anymore; if required, Hanning smoothing can be applied during post-processing (e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>].). The Gibbs effect can also be reduced by using higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reducing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on-source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and number of polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The Proposal Submission Tool (PST) == The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. <!-- link to Michaels docs--> [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used for a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz for full, 1MHz for dual and 0.5MHz for single polarization products. No narrow subbands can be chosen. == Setting up a Spectral Observation using the Observation Preparation Tool (OPT) == The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the [http://evlaguides.nrao.edu/index.php?title=Category:OPT-QuickStart OPT QuickStart Guide], and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. === Bandpass Setup === All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within an SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as RFI lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> === Phase/Complex Gain Calibration === The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. === Correlator Resources Setup === '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line observations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> ==== Using Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> e5bcd93b1169a3b6cd15bf5fc9d388d246d45bb1 1323 1322 2012-07-02T06:06:21Z Jott 8 /* Current->Revised OSS Guidelines */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> = Introducing VLA Spectral Line Observing = The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization = Line Frequencies and Velocity Conventions = == Line Rest Frequencies == Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. = The WIDAR Correlator = == Basebands == Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Baseband Tuning Restrictions === The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation == Fixed 128MHz Subbands and 128MHz "Suckouts" == After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] == Correlator Resources and Subband Placement == Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. === Narrow Subbands with the 8-bit sampler === Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> === Standard Subbands === Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> === Recirculation === "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. == Narrow Subbands with the 3-bit sampler == As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. == Data Rate Limits == Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 20MB/s for regular and more for shared risk observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary for details]. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. = Tips for Planning, Setup, and Processing of VLA Spectral Line Observations = '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization == Considerations for Planning Subband Bandwidths and Resolution == === Bandwidths required for UV continuum subtraction === When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. === Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon, a sinc function that zig-zags on alternating channels. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing is the most effective method, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the pre-upgrade VLA days, the correlator design had a realtively short, truncated lag spectrum, which could result in prominent Gibbs ringing. To avoid this effect, Hanning smoothing was frequently applied online during the observations. With the new WIDAR capabilities of the upgraded VLA, however, ringing is very rare and only observed for extremely strong maser or RFI sources. Consequently, the VLA does not support online Hanning smoothing anymore; if required, Hanning smoothing can be applied during post-processing (e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>].). The Gibbs effect can also be reduced by using higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reducing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on-source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and number of polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The Proposal Submission Tool (PST) == The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. <!-- link to Michaels docs--> [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used for a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz for full, 1MHz for dual and 0.5MHz for single polarization products. No narrow subbands can be chosen. == Setting up a Spectral Observation using the Observation Preparation Tool (OPT) == The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the [http://evlaguides.nrao.edu/index.php?title=Category:OPT-QuickStart OPT QuickStart Guide], and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. === Bandpass Setup === All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within an SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as RFI lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> === Phase/Complex Gain Calibration === The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. === Correlator Resources Setup === '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line observations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> ==== Using Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. The new Karl G. Jansky Very Large Array (VLA) correlator is extremely powerful in its spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from from 2MHz down to the the Hz regime. Here is a guide to access that spectral line power and we describe what is offered for the August 2012 deadline. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> == VLA Spectral Line Observing == This guide is to help understand and set up spectral line observations at the VLA. The new, wide bandwidths of the VLA allow users to observe up to 8GHz of spectral bandwidth at a time. Apart from extreme continuum sensitivity, the wide bands of the VLA can be used to observe multiple spectral lines simultaneously. Furthermore, the WIDAR correlator is extremely flexible and acts, fundamentally, like up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA can * deliver continuous spectral coverage up to a full width of 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place correlator subbands on them; this enables observations of multiple spectral lines at once * use up to 64 subbands at a time that are independently tunable, and can be configured in different spectral bandwidths, channel numbers, and number of polarization products * derive the frequency from the velocity of a given spectral line (Doppler Setting) * Dynamically schedule the observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution == Line Rest Frequencies == There are a number of online tools available that help spectral line observers. The recommended line rest frequency catalog for VLA and ALMA is [http://splatalogue.net Splatalogue] which contains data from the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines and data from other resources. == Observing Frequency and Velocity Definitions == The first step is to determine the observing frequency <math>\nu</math> of the spectral line. This is derived from the radial velocity <math>v</math> of the source and the rest frequency <math>\nu_0</math> of the spectral line. A full relativistic calculation shows that the velocity <math>v</math> is determined via <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math> and is called the ''relativistic velocity''. The equation is a bit cumbersome to use and in astronomy two different approximations are typically used instead: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. At low velocities that difference is small but <math> v^{optical}</math> and <math> v^{radio}</math> diverge more and more for large values. Traditionally, the optical velocities are predominantly used for extragalactic and the radio velocities for Galactic targets. At significant redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math> where the redshifted <math>\nu</math> can now be used as the input rest frequency for the observations. The velocities that are derived based on such a redshifted rest frequency will be correctly scaled for the spread of the velocity scale that is caused by the redshift. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. If one measures a velocity it is therefore necessary to correct for such motions and define the frame in which the velocities are measured. There are various rest frames used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the the following three reference frames is commonly used: * '''Topocentric''', the natural velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''': the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise LSR naming is usually used synonymously with more modern LSRK definition. * '''Barycentric''': a common frame that has virtually replaced the older heliocentric standard. Given the small difference between barycentric and heliocentric, they were frequently used interchangeably. A full list of CASA supported reference frames is provided in the [[http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook]] and also on the [[http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage]] == Doppler Correction == A telescope frequently operates at a fixed sky frequency (Topocentric velocity frame). Any spectral line will thus shift during any observation campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position on the sky (see above). Observing campaigns that span over a year, may see spectral lines to shift in frequency by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Side note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. There are different ways to do this: * use the same sky frequency for all observations. The shift of the line (maximum of <math>\pm</math>30km/s) is accommodated by a wide bandwidth that covers the line and its width at any time of the observation campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate a fixed sky frequency at the beginning of an observation, then keep it fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then minimized to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected by post processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The post-processing regridding of the line in CASA can be either done directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. == Gibbs Phenomenon and Hanning smoothing == For very sharp spectral or lag features, a Fourier transform can prominently display a sinc function, a channel by channel fluctuation of the amplitude. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. The smoothing kernel to be used is a Hanning smoothing function which sports a triangular kernel with the central channel being weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the VLA days, the correlator design (with a truncated lag spectrum) could show the Gibbs phenomenon relatively prominently and Hanning smoothing was frequently applied online during the observations to accommodate for the effect and to save disk space. The Gibbs phenomenon is much less common for the VLA due to a better correlator design of WIDAR. Only very strong maser or RFI sources may exhibit the typical "ringing" feature of the Gibbs phenomenon. In addition, data can be stored rather cheaply so there is no need for data size reduction via online Hanning smoothing anymore. Consequently, the VLA does not support online Hanning smoothing. If Hanning smoothing is required, it has to be performed in post-processing, e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>]. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] The sensitivity of spectral line data is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to enter sensitivity limits and provides the required time on source given a frequency, weather, weighting scheme, polarization products, and bandwidth of the observations. The '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The WIDAR Correlator == Together with the wide, instantaneous receiver bandwidths, the WIDAR correlator at the VLA is very flexible and provides a number of setup options that are relevant for spectral line observing. === Basebands === Let's start with the basics: A signal from the telescope enters the WIDAR and along that path, it is passing analog filters that define the <b>basebands</b>. The basebands are actually baseband pairs to cover the L and R polarizations. They are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. In the 8-bit mode, the samplers deliver two independably tuneable basebands (dubbed AC0 and BD0) with 1 GHz bandwidth each. Using the 3-bit samplers, there are 4 baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. Each baseband pair can be set to a baseband sky frequency. Alternatively, the sky frequency can be calculated using the Doppler Setting tool in OPT. When the signal enters WIDAR, the basebands pass through a number of digital filters, each 128MHz wide, the fixed 128MHz subbands. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Fixed 128MHz Subbands and 128MHz "Suckouts" === After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. As each fixed 128MHz subband has some filter shape with soft corners, the sensitivity of the VLA drops to about half its value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are a few options to account and interpolate over the suckouts: * Easiest method is to avoid them in your spectral setup. Try to set the baseband frequency in a way such that any interesting lines do not fall in the suckouts. We offer a [https://webtest.aoc.nrao.edu/cgi-internal/nroy/work.py spectral line setup tool] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. * Observe with two basebands shifted by 10-64MHz apart. This will ensure that at one of the baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. To obtain more spectral channels, one can consider to using single polarizations on the basebands. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] === Narrow Subbands === For every baseband, there can be a maximum of 64 subbands (for the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations). These subbands can be one of the fixed 128MHz wide subbands or more narrow, tunable subbands. Fundamentally, each narrow subband can be between 128 MHz (corresponding to a fixed subband) and 31.25 kHz wide and contains 64 channels when all four RR, LL, RL, LR polarization products are required (full polarization), 128 channels in dual polarization mode (RR & LL), and 256 channels for single, RR or LL polarization products. The two tables display the options for full and dual polarizations with the corresponding velocity widths and channelizations. Higher spectral resolution can be achieved with baseline board stacking and recirculation (August 2012: only shared risk) as described below. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> ==== WIDAR Tuning Restrictions ==== The baseband (pairs) cannot be entirely independently tuned. The following restrictions apply: * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * In Ka band, 3-bit samplers, only one baseband can be below 32GHz and that must be BD * The 8-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation For the narrow subbands: * Narrow subbands can not go across a 128MHz suckout frequency. Any narrow subband needs to be entirely within a 128MHz fixed subband. === Correlator Baselineboards === Correlator baselineboards (BlBs, also named "BL.BPS" in OPT) are independent units that can be used for separate subbands. WIDAR has 64 BlB pairs (for the polarizations) and thus supports a maximum of 64 subbands. === Baselineboard Stacking === If not all of the 64 subbands are used, the remaining BlBs can be used to obtain more channels per subband. This method is called "baselineboard stacking" and each additional BlB for a subband adds another 64 channels in full and 128 channels in dual polarization modes. In OPT's resource tool (RCT), this can be set via the BLB.BPS dropdown menu that is available for each subband. E.g. the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz subbands, each with 128 channels dual polarization as in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baseline board stacking is also extremely useful as it allows the use of wider bandwidths for each subband while still maintaining a high number of channels. This minimizes the number of filter edges that are unavoidable for each subband. === Recirculation === Another way to obtain more spectral channels for a given subband is called "recirculation". Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The basic correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. This setup, however, is only available for shared risk observations for the August 2012 deadline. === Data Rate Limits === Baselineboard stacking, recirculation, and time resolution, however, can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 15MB/s for regular and higher data rates for shared risk observing. The OPT instrument configuration calculates data rates based on the spectral line setup and the limit of 60MB/s should not be exceeded for any observational setup. == Planning and Setup of VLA Spectral Line Observations== '''Reminder:''' For the August 1 deadline, the following capabilities are offered (more options are available for "shared risk" observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]): * two polarization pairs 1GHz each (using the 8-bit samplers) * up to 16 subband polarization pairs per baseband polarization pair * independent tuning of all subbands * independent subband bandwidth in the 31.5-128MHz range * independent number of channels, up to 2048 channels for single polarization product, 1024 dual, 512 full polarization; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * Doppler setting * The maximum data rate, however, cannot exceed 15MB/s * for the wideband 3-bit samplers, 2x2GHz baseband polarization pairs can be set up with 2MHz resolution === The Proposal Submission Tool (PST) === To comply with all the restrictions listed above for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. The PST is accessible via [http://my.nrao.edu my.nrao.edu]. You may need to register at if you do not yet have an account. [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] The 8-bit sampler data can be set up with the flexibility mentioned above. The figure to the right shows an example where 9 subbands were chosen in Ka band, four in the first and 5 in the second baseband. The setup features different subband bandwidths, polarization products and baselineboard stacking (up to 16 for a couple of subbands). A total of 49 baseline boards are used for this configuration. A full description of the tool '''is provided here --- ??? link to Michael's manual'''. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz and no narrow subbands can be chosen. === Setting up a Spectral Observation using the Observation Preparation Tool (OPT) === The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for the VLA. An SB is a block of time during which an observation is executed. A full project may constitute of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Please also read the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] that contains the latest information on the OPT as well as a comprehensive OPT manual. ==== Frequency Setup ==== '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' One of the most important parts of the setup is to chose the appropriate frequency setup for WIDAR. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk ==== Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation, based on the time of the observation, the velocity, position and rest frequency of a source and line. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. ==== Bandpass Setup ==== All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and in many cases can double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be good enough for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. <!-- Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. Attenuators are typically set up at the dummy scan for each frequency setup at the start of the SB and will be fixed for the full execution of the SB. In the case that this "set-and-remember" mechanism fails, it is possible to track down such attenuator changes in the data using the switched power information; the On - Off power ('PDIF' in AIPS) will show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation. If there is only one observation of the bandpass calibrator, --> The observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent. But we advise to use that option only when absolutely necessary. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as rfi lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. ==== Phase/Complex Gain Calibration ==== The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should stick with the correlator setup for the target and the complex gain calibrator. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. ==== Continuum Subtraction ==== To allow for good continuum subtraction, it is important that enough line-free channels are being observed on either frequency end of the spectral line. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. <!-- The continuum of an observation can be either subtracted in the ''uv'' plane ([http://casa.nrao.edu/docs/TaskRef/uvcontsub-task.html ''uvcontsub''] in CASA) or in the image domain ([http://casa.nrao.edu/docs/TaskRef/imcontsub-task.html ''imcontsub'']). In most cases, the uv-domain is preferred but for high dynamic range imaging, or imaging close or beyond the edge of the FWHM of the primary beam, it may be better to subtract the continuum using line-free channels in the image cube, or to Fourier transform the com=ntinuum image and subtract these data from the visibilities ([http://casa.nrao.edu/docs/TaskRef/uvsub-task.html ''uvsub'']). All methods have in common that the continuum can only be determined accurately when enough line-free channels are observed that are necessary to derive a good model for the continuum. If the continuum sources exhibit significant spectral slope or even curvature, it is advisable to go to even larger bandwidth covering more line-free frequency space and to sample the continuum with many channels to ensure that a reasonable higher-order fit can be derived and subtracted from the line+continuum data. --> ==== High Dynamic Range Imaging ==== For very strong and narrow spectral features (typically thousands of Jy strong sources), one may see the Gibbs phenomenon (ringing) and Hanning smoothing may need to be applied (see above). This needs to be done in post-processing. The effect, however, can be reduced by using a higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reduceing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. <!-- = Spectral Line Observing = ''' everything below this section may be obsolete''' == Current->Revised OSS Guidelines == * An Overview of the VLA (last paragraph) The VLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for VLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. It is important to realise that the VLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the VLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the pre-upgrade VLA are strongly advised to consult Correlator Configurations. * Limitations on Imaging Performance (Sidelobes from Strong Sources) An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. * Correlator Configurations All observations with the VLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the VLA correlator. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. --> 0281653a859f8c0d5335101da596a4eec3794a60 1324 1323 2012-07-02T06:07:33Z Jott 8 wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> = Introducing VLA Spectral Line Observing = The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization = Line Frequencies and Velocity Conventions = == Line Rest Frequencies == Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. = The WIDAR Correlator = == Basebands == Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed AC0 and BD0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1C1, A2C2, B1D1, B2D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Baseband Tuning Restrictions === The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be BD0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation == Fixed 128MHz Subbands and 128MHz "Suckouts" == After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] == Correlator Resources and Subband Placement == Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. === Narrow Subbands with the 8-bit sampler === Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> === Standard Subbands === Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> === Recirculation === "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 deadline. == Narrow Subbands with the 3-bit sampler == As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. == Data Rate Limits == Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 20MB/s for regular and more for shared risk observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary for details]. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. = Tips for Planning, Setup, and Processing of VLA Spectral Line Observations = '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization == Considerations for Planning Subband Bandwidths and Resolution == === Bandwidths required for UV continuum subtraction === When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. === Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon, a sinc function that zig-zags on alternating channels. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing is the most effective method, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the pre-upgrade VLA days, the correlator design had a realtively short, truncated lag spectrum, which could result in prominent Gibbs ringing. To avoid this effect, Hanning smoothing was frequently applied online during the observations. With the new WIDAR capabilities of the upgraded VLA, however, ringing is very rare and only observed for extremely strong maser or RFI sources. Consequently, the VLA does not support online Hanning smoothing anymore; if required, Hanning smoothing can be applied during post-processing (e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>].). The Gibbs effect can also be reduced by using higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reducing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on-source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and number of polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The Proposal Submission Tool (PST) == The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. <!-- link to Michaels docs--> [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used for a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz for full, 1MHz for dual and 0.5MHz for single polarization products. No narrow subbands can be chosen. == Setting up a Spectral Observation using the Observation Preparation Tool (OPT) == The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the [http://evlaguides.nrao.edu/index.php?title=Category:OPT-QuickStart OPT QuickStart Guide], and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. === Bandpass Setup === All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within an SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as RFI lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> === Phase/Complex Gain Calibration === The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. === Correlator Resources Setup === '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line observations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> ==== Using Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once for each baseband at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. 6c789729841f96a3172656c0de86fbcb127aa0d7 1325 1324 2012-07-02T06:18:29Z Jott 8 wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> = Introducing VLA Spectral Line Observing = The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization = Line Frequencies and Velocity Conventions = == Line Rest Frequencies == Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (<math> v </math> and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT '''for each baseband''' (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. = The WIDAR Correlator = == Basebands == Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed A0/C0 and B0/D0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1/C1, A2/C2, B1/D1, B2/D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Baseband Tuning Restrictions === The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be B0/D0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation == Fixed 128MHz Subbands and 128MHz "Suckouts" == After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] == Correlator Resources and Subband Placement == Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. === Narrow Subbands with the 8-bit sampler === Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> === Standard Subbands === Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> === Recirculation === "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 proposal deadline. == Narrow Subbands with the 3-bit sampler == As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. == Data Rate Limits == Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 20MB/s for regular and more for shared risk observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary for details]. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. = Tips for Planning, Setup, and Processing of VLA Spectral Line Observations = '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 proposal deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization == Considerations for Planning Subband Bandwidths and Resolution == === Bandwidths required for UV continuum subtraction === When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. === Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon, a sinc function that zig-zags on alternating channels. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing is the most effective method, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the pre-upgrade VLA days, the correlator design had a realtively short, truncated lag spectrum, which could result in prominent Gibbs ringing. To avoid this effect, Hanning smoothing was frequently applied online during the observations. With the new WIDAR capabilities of the upgraded VLA, however, ringing is very rare and only observed for extremely strong maser or RFI sources. Consequently, the VLA does not support online Hanning smoothing anymore; if required, Hanning smoothing can be applied during post-processing (e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>].). The Gibbs effect can also be reduced by using higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reducing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on-source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and number of polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The Proposal Submission Tool (PST) == The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. <!-- link to Michaels docs--> [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used for a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz for full, 1MHz for dual and 0.5MHz for single polarization products. No narrow subbands can be chosen. == Setting up a Spectral Observation using the Observation Preparation Tool (OPT) == The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the [http://evlaguides.nrao.edu/index.php?title=Category:OPT-QuickStart OPT QuickStart Guide], and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. === Bandpass Setup === All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within an SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as RFI lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> === Phase/Complex Gain Calibration === The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. === Correlator Resources Setup === '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line observations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> ==== Using Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once '''for each baseband''' at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. d75f8226af35d76e35ef6d4f9a3d417777db6795 1326 1325 2012-07-02T06:19:57Z Jott 8 /* Observing Frequency and Velocity Definitions */ wikitext text/x-wiki The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> = Introducing VLA Spectral Line Observing = The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization = Line Frequencies and Velocity Conventions = == Line Rest Frequencies == Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (v) and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT '''for each baseband''' (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. = The WIDAR Correlator = == Basebands == Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed A0/C0 and B0/D0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1/C1, A2/C2, B1/D1, B2/D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Baseband Tuning Restrictions === The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be B0/D0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation == Fixed 128MHz Subbands and 128MHz "Suckouts" == After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] == Correlator Resources and Subband Placement == Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. === Narrow Subbands with the 8-bit sampler === Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> === Standard Subbands === Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> === Recirculation === "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 proposal deadline. == Narrow Subbands with the 3-bit sampler == As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. == Data Rate Limits == Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 20MB/s for regular and more for shared risk observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary for details]. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. = Tips for Planning, Setup, and Processing of VLA Spectral Line Observations = '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 proposal deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization == Considerations for Planning Subband Bandwidths and Resolution == === Bandwidths required for UV continuum subtraction === When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. === Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon, a sinc function that zig-zags on alternating channels. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing is the most effective method, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the pre-upgrade VLA days, the correlator design had a realtively short, truncated lag spectrum, which could result in prominent Gibbs ringing. To avoid this effect, Hanning smoothing was frequently applied online during the observations. With the new WIDAR capabilities of the upgraded VLA, however, ringing is very rare and only observed for extremely strong maser or RFI sources. Consequently, the VLA does not support online Hanning smoothing anymore; if required, Hanning smoothing can be applied during post-processing (e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>].). The Gibbs effect can also be reduced by using higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reducing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on-source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and number of polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The Proposal Submission Tool (PST) == The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. <!-- link to Michaels docs--> [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used for a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz for full, 1MHz for dual and 0.5MHz for single polarization products. No narrow subbands can be chosen. == Setting up a Spectral Observation using the Observation Preparation Tool (OPT) == The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the [http://evlaguides.nrao.edu/index.php?title=Category:OPT-QuickStart OPT QuickStart Guide], and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. === Bandpass Setup === All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within an SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as RFI lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> === Phase/Complex Gain Calibration === The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. === Correlator Resources Setup === '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line observations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> ==== Using Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once '''for each baseband''' at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. 617e49f02fbac791407153fffb09eea4f14d9362 1327 1326 2012-07-02T14:05:30Z Jott 8 wikitext text/x-wiki <math>f</math> The new Karl G. Jansky Very Large Array (VLA) has extremely powerful and versatile spectral capabilities. The final specifications include up to 4 million channels that can be distributed in up to 64 subbands with a spectral resolution from 2MHz down to the the Hz regime. This guide is intended to clarify the spectral line capabilities of the VLA, with a focus on capabilities offered for the August 2012 VLA proposal deadline, and enable users to plan, prepare, and process spectral observations. <!-- [http://www.vla.nrao.edu/astro/guides/sline/current/ VLA spectral line guide] Contents INTRODUCTION SYSTEM SPECIFICATIONS Receivers and IF System The Local Oscillator Chain The Correlator Total Number of Channels Some Advanced Spectral Line Topics Creative Use Of The Spectral Line System High Accuracy Spectral Line Polarization Observations The Lag Spectrum, Gibbs Phenomenon and Hanning Smoothing Continuum Observations in Line Mode OBSERVATIONAL CONSIDERATIONS Amplitude and Phase Calibration Bandpass Calibration Continuum Subtraction Interference Bandwidth and Time Smearing Determining the Observing Frequency Doppler tracking and dopset Velocity Definition Summing Velocities Velocity Rest Frame Running dopset Setting the LO Chain: loser System Temperature Corrections Changes In T$_{sys}$ With Elevation (T$_{spill}$) Contributions to T$_{sys}$ From Strong Lines REFERENCES Bandwidth and Number of Channels Normal Mode On-line Hanning Smoothing Option --> = Introducing VLA Spectral Line Observing = The newly available wide bandwidths of the VLA allow users to observe up to 8GHz of bandwidth at a time. All observations with the upgraded VLA are inherently spectral observations, including those intended for continuum science. The VLA's improved sensitivity and wide bandwidths greatly enhance the VLA's functionality for spectral line purposes, enabling simultaneous imaging of multiple spectral lines. The WIDAR correlator is extremely flexible and can act as up to 64 independent correlators with different bandwidths, channel numbers, polarization products, and observing frequencies. The final VLA will be able to * deliver continuous spectral coverage of up to 8GHz * access 1GHz or 2GHz chunks in each receiver band (called basebands) and place multiple correlator subbands within them * place up to 64 independently tunable subbands within a baseband; these can be configured with different bandwidths, channel numbers, and polarization products * tune the baseband and subband frequencies according to the object's velocity with respect to the earth (Doppler Setting) * dynamically schedule observations to use the best weather conditions for high frequency, high scientific impact projects <!-- * Post-processing: The [http://casa.nrao.edu CASA] package is the main data reduction software for VLA and ALMA and contains cutting edge data reduction code with continuum and spectral line processing being the main focus. The software and reference guides can be obtained on the [http://casa.nrao.edu CASA homepage]. The [http://casaguides.nrao.edu CASAguides wiki] contains guides on VLA spectral line data reduction as well as some hints, tips and tricks on using CASA and the visualization tools that are designed to display spectral data cubes --> The following capabilities are offered for standard observing in the August 2012 proposal submission deadline. Greater flexibility is available through "shared risk" observing, as discussed below and detailed in the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization = Line Frequencies and Velocity Conventions = == Line Rest Frequencies == Spectral line calalogues available online are useful for selecting targeted line rest frequencies. The recommended catalog for VLA and ALMA observing is [http://splatalogue.net Splatalogue] which contains molecular line data from sources including the [http://physics.nist.gov/cgi-bin/micro/table5/start.pl Lovas catalog], the [http://spec.jpl.nasa.gov/ JPL/NASA molecular database]], the [http://www.astro.uni-koeln.de/cdms/ Cologne Database for Molecular Spectroscopy], as well as radio recombination lines. == Observing Frequency and Velocity Definitions == The frequency at which we must tune the correlator in order to observe a spectral line (<math> \nu </math>) is derived from the radial velocity of the source (v) and the rest frequency of the spectral line (<math> \nu_0 </math>). The ''relativistic velocity'', or true radial velocity, is related to the observed and rest frequencies by <math> v = \frac{\nu_0^{2} - \nu^{2}}{\nu_0^{2}+\nu^{2}}</math>. This equation is a bit cumbersome to use; in astronomy two different approximations are typically used: * '''Optical Velocity''' <math> v^{optical} = \frac{\lambda-\lambda_0}{\lambda_0}\,\,c = cz </math> (<math>z</math> is the redshift of the source) * '''Radio Velocity''' <math> v^{radio} = \frac{\nu_0-\nu}{\nu_0}\,\,c = \frac{\lambda-\lambda_0}{\lambda}\,\,c \neq v^{optical}</math> The radio and optical velocities are not identical. Particularly,<math> v^{optical}</math> and <math> v^{radio}</math> diverge for large velocities. Optical velocities are predominantly used for extragalactic and radio velocities for Galactic targets. At high redshifts, it is advisable to place the zero point of the velocity frame into the source via <math> \nu = \frac{\nu_0}{z+1} </math>, where the redshifted <math>\nu</math> can now be used as the input frequency for the observations. This method will appropriately apply the redshift correction to the channel and line widths and the resulting velocities are also intrinsic to the source. == Velocity Frames == The earth rotates, revolves around the sun, rotates around the galaxy, moves within the Local Group, and shows motion against the cosmic microwave background. Any source velocity must therefore always be specified relative to a reference frame. Various velocity rest frames are used in the literature. The following table lists their name, the motion that is corrected for, and the maximum amplitude of the velocity correction. Each rest frame correction is incremental to the preceding row: <TABLE CELLPADDING=4 BORDER="1"> <tr><td><b>Rest Frame Name</b></td><td><b>Rest Frame</b></td><td><b>Correct for</b></td><td><b>Max amplitude [km/s]</b></td></tr> <tr><td>Topocentric</td><td>Telescope</td><td>Nothing</td><td>0</td></tr> <tr><td>Geocentric</td><td>Earth Center</td><td>Earth rotation</td><td>0.5</td></tr> <tr><td>Earth-Moon Barycentric</td><td>Earth+Moon center of mass</td><td>Motion around Earth+Moon center of mass</td><td>0.013</td></tr> <tr><td>Heliocentric</td><td>Center of the Sun</td><td>Earth orbital motion</td><td>30</td></tr> <tr><td>Barycentric</td><td>Earth+Sun center of mass</td><td>Earth+Sun center of mass</td><td>0.012</td></tr> <tr><td>Local Standard of Rest (LSR)</td><td>Center of Mass of local stars</td><td>Solar motion relative to nearby stars</td><td>20</td></tr> <tr><td>Galactocentric</td><td>Center of Milky Way</td><td>Milky Way Rotation</td><td>230</td></tr></td></tr> <tr><td>Local Group Barycentric</td><td>Local Group center of mass</td><td>Milky Way Motion</td><td>100</td></tr> <tr><td>Virgocentric</td><td>Center of the Local Virgo supercluster</td><td>Local Group motion</td><td>300</td></tr> <tr><td>Cosmic Microwave Background</td><td>CMB</td><td>Local Supercluster Motion</td><td>600</td></tr> </table> <p> The velocity frame should be chosen based on the science. For most observations, however, one of the following three reference frames is commonly used: <br> * '''Topocentric''' is the velocity frame of the observatory (defining the sky frequency of the observations). Visibilities in a measurement set are typically stored in this velocity frame. * '''Local Standard of Rest''' is the native output of images in CASA. Note that there are two varieties of LSR: the kinematic LSR ('''LSRK''') and the dynamic (LSRD) definitions for the kinematic and dynamic centers, respectively. In almost all cases LSRK is being used and the less precise name 'LSR' is usually used synonymously with more modern LSRK definition. * '''Barycentric''' is a commonly used frame that has virtually replaced the older heliocentric standard. Given the small difference between the barycentric and heliocentric frames, they were frequently used interchangeably. <br> A full list of CASA supported reference frames is provided in the [http://casa.nrao.edu/docs/userman/UserMan.html CASA reference Manual and Cookbook] and also on the [http://casaguides.nrao.edu/index.php?title=Velocity_Reference_Frames casaguides.nrao.edu webpage] == Doppler Correction == A telescope naturally operates at a fixed sky frequency (Topocentric velocity frame) which can be adjusted to account for the motion of the earth. A spectral line's observed frequency will shift during any observing campaign. Within a day, the rotation of the earth dominates and the line may shift up to <math>\pm</math>0.5km/s, depending on the position of the source on the sky (see above). Observing campaigns that span a year may have spectral lines that shift by up to <math>\pm</math>30km/s due to the earth's motion around the sun. <i>Note: As a rule of thumb, 1 MHz in frequency corresponds roughly to <math>x</math> km/s for the line at a wavelength of <math>x</math> in mm. E.g., at a wavelength of 7mm, 1MHz corresponds to about 7km/s in velocity, at 21cm 1MHz corresponds roughly to 210km/s. </i> Using this rule of thumb, a line may shift by up to <math>\pm</math>5MHz in Q-band and by up to <math>\pm</math>0.15MHz in L-band over the course of a year. This needs to be taken into account when setting up the observations. This issue can be handled in different ways: * use the same sky frequency for all observations, accommodating the line shift (maximum of <math>\pm</math>30km/s) by using a wide enough bandwidth to cover the line at any time in the observing campaign. The data is later regridded in CASA to a common LSRK or BARY velocity frame. The sky frequency of an observation can be computed with the [http://www.vla.nrao.edu/astro/guides/dopset/ Dopset tool] for a given time. One may find the LST dates for an observation on the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi VLA Schedule Page]. * calculate the sky frequency at the beginning of an observing block and keep this fixed for the duration of the scheduling block. This is called '''Doppler Setting''' and offered by OPT '''for each baseband''' (currently only in OSRO mode). The line shift is then reduced to the rotation of the earth (maximum amplitude <math>\pm</math>0.5km/s). This small shift is corrected in data processing. * change the sky frequency continuously to keep the line at the same position in the band. This method is called '''Doppler tracking''' and was standard for the pre-upgrade VLA. <i>The new VLA does NOT support Doppler tracking.</i> The WIDAR correlator offers enough bandwidth and spectral channels to cover any line shift and post-processing regridding needs. In addition, a non-variable sky frequency may also deliver a more robust calibration and overall system stability. The regridding of the spectrum can be completed during data processing in CASA, either directly during imaging in the task <tt>clean</tt>, or alternatively with the task <tt>cvel</tt>. The regridding works well when the spectral features are sampled with at least 4 channels. = The WIDAR Correlator = == Basebands == Let's start with the basics: A signal from the telescope enters WIDAR, is split into its left and right hand circular polarizations, and passes through analog filters that define the <b>basebands</b>. Basebands set the spectral range that can be accessed by subbands, and they come in baseband pairs to cover L and R polarizations. Each baseband pair can be set to one baseband sky frequency. Basebands are the most fundamental spectral ranges delivered from the samplers and digitizers to the correlator. With 8-bit sampling, the samplers deliver two independently tuneable baseband pairs (dubbed A0/C0 and B0/D0) with 1 GHz bandwidth each. 3-bit sampling provides four baseband pairs (A1/C1, A2/C2, B1/D1, B2/D2), each of them 2GHz wide. [[File:WIDARcorrelatorbands.png|300px|thumb|right|Baseband with WIDAR subbands]] === Baseband Tuning Restrictions === The following restrictions apply to baseband tuning: * With 8-bit sampling in Ka band, only one baseband can be below 32GHz and that must be B0/D0 * 3-bit samplers can only be used in C-band or above, where the instantaneous frequency width of the receiver is larger than 2GHz. * The 3-bit A1C1 baseband frequency can have a separation of max. 4GHz from the A2D2 baseband. B1D1 and B2D2 have the same restriction of 4GHz maximum separation == Fixed 128MHz Subbands and 128MHz "Suckouts" == After filtering through the basebands, the signal enters the correlator and is split into fixed, 128MHz wide subbands. They are placed adjacently to cover the full width of the basebands. Narrow subbands can be arranged within these fixed 128MHz subbands. Because each fixed 128MHz subband has a filter shape with soft corners, the sensitivity of the VLA drops to about half its maximum value between any two fixed 128MHz subbands. These frequency ranges are called "128 MHz Suckouts". There are two primary options for dealing with the suckouts: 1. Try to set the baseband frequency such that targeted lines do not fall in the suckouts. We offer the [https://e2e.nrao.edu/tune.shtml spectral line setup tool "TUNE"] that can be used to maximize the frequency separation between a number of spectral lines and the suckouts. 2. If it is not possible to obtain coverage of all of your lines using the above method, observe with two basebands shifted by 10-64MHz apart. This will ensure that at one baseband covers the suckouts of the other baseband with full sensitivity. An example is given in the figures to the right. [[File:BlankFieldRMS.AC.png|300px|thumb|right| rms noise in a blank field as a function of frequency for one baseband consisting of 8 contiguous sub-bands. Note the increased rms noise at the subband edges.]] [[File:BlankFieldRMS.interlace.png|300px|thumb|right| rms noise in a blank field as a function of frequency for two basebands consisting of 8 contiguous sub-bands, where the basebands are separated by one-half of the subband width. Wherever signal in one baseband is compromised by edge effects, data from the other subband are substituted.]] == Correlator Resources and Subband Placement == Correlator baselineboards (BlBs, also named "BL.BPS" for "baselineboard pairs" in the OPT) are independent hardware units that are allocated to narrow subbands. <b>WIDAR has 64 baselineboard pairs.</b> As a result, WIDAR supports a maximum of 64 subband pairs (again, pairs to cover the two polarizations), but the number of subbands depends on the subband setups as described below. For the same number of channels, single (RR or LL), dual (RR & LL), and full (RR, LL, RL, and LR) polarization products require 1, 2, and 4 times the correlator baselineboards respectively. Similarly, doubling the number of channels in a subband doubles the number of correlator resources used. === Narrow Subbands with the 8-bit sampler === Narrow subbands define the frequency ranges in which the spectrum is measured. Narrow subbands with bandwidths between 31.25kHz and 128MHz can be arranged to obtain desired frequency coverage and spectral resolution within the baseband. Each narrow subband needs to be entirely within a fixed 128MHz subband, as <b> narrow subbands cannot cross a 128MHz suckout frequency.</b> The baseband center frequency can be shifted and subband bandwidths and frequencies must be arranged to avoid the suckouts. This implies that the 128MHz fixed subbands cannot be moved as they would fall on a suckout at any frequency offset from a 128MHz "raster" within the baseband. All subbands less than 128MHz in width, however, and can be independently tuned as long as they do not cross a suckout. Furthermore, <i>all subbands can be set up with different bandwidths, frequency resolutions, channel numbers, and polarization products.</i> === Standard Subbands === Standard subbands allocate a single baselineboard pair to a each subband. Standard subbands contain 64 channels when full polarization (RR,LL, RL & LR) products are required, 128 channels in dual polarization mode (RR & LL), and 256 channels for single polarization (RR or LL). Options for full and dual polarization subbands, with frequency and velocity resolutions, are shown in the following tables. For the August 2012 deadline, we offer up to 16 subbands per baseband for normal observations, and up to 64 for shared risk observations. * Full polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage per sub-band (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>64</td> <td>2000</td> <td>600/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>64</td> <td>1000</td> <td>300</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>64</td> <td>500</td> <td>150</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>64</td> <td>250</td> <td>75</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>64</td> <td>125</td> <td>37.5</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>64</td> <td>62.5</td> <td>19</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>64</td> <td>31.25</td> <td>9.4</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>64</td> <td>15.625</td> <td>4.7</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>64</td> <td>7.813</td> <td>2.3</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>64</td> <td>3.906</td> <td>1.2</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>64</td> <td>1.953</td> <td>0.59</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>64</td> <td>0.977</td> <td>0.29</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>64</td> <td>0.488</td> <td>0.15</td> <td>9.375</td> </tr> </table> * Dual Polarization <table class="grid listing" border="1"> <tr valign="top"> <th>Sub-band BW (MHz) </th><th>Number of channels/poln product </th><th>Channel width (kHz) </th><th>Channel width (km/s at 1 GHz) </th><th>Total velocity coverage (km/s at 1 GHz) </th> </tr> <tr valign="top"> <td>128</td> <td>128</td> <td>1000</td> <td>300/ν(GHz)</td> <td>38,400/ν(GHz)</td> </tr> <tr valign="top"> <td>64</td> <td>128</td> <td>500</td> <td>150</td> <td>19,200</td> </tr> <tr valign="top"> <td>32</td> <td>128</td> <td>250</td> <td>75</td> <td>9,600</td> </tr> <tr valign="top"> <td>16</td> <td>128</td> <td>125</td> <td>37.5</td> <td>4,800</td> </tr> <tr valign="top"> <td>8</td> <td>128</td> <td>62.5</td> <td>19</td> <td>2,400</td> </tr> <tr valign="top"> <td>4</td> <td>128</td> <td>31.25</td> <td>9.4</td> <td>1,200</td> </tr> <tr valign="top"> <td>2</td> <td>128</td> <td>15.625</td> <td>4.7</td> <td>600</td> </tr> <tr valign="top"> <td>1</td> <td>128</td> <td>7.813</td> <td>2.3</td> <td>300</td> </tr> <tr valign="top"> <td>0.5</td> <td>128</td> <td>3.906</td> <td>1.2</td> <td>150</td> </tr> <tr valign="top"> <td>0.25</td> <td>128</td> <td>1.953</td> <td>0.59</td> <td>75</td> </tr> <tr valign="top"> <td>0.125</td> <td>128</td> <td>0.977</td> <td>0.29</td> <td>37.5</td> </tr> <tr valign="top"> <td>0.0625</td> <td>128</td> <td>0.488</td> <td>0.15</td> <td>18.75</td> </tr> <tr valign="top"> <td>0.03125</td> <td>128</td> <td>0.244</td> <td>0.073</td> <td>9.375</td> </tr> </table> === Baselineboard Stacking === In order to obtain a larger number of channels per subband, a method called "baselineboard stacking" can be used. Baselineboard stacking allows a larger number of baseline boards to be allocated to a single subband, increasing the number of channels within the subband. To double the number of channels in a subband, simply double the number of baselineboards allocated to the subband. Thus each additional BlB for a subband adds another 64 channels in full and 128 channels in dual, and 256 channels in single polarization modes. As an example, the full 2GHz bandwidth of the 8-bit samplers can be covered by 16 128MHz standard narrow subbands, each with 128 channels in dual polarization, as shown in the table above. The 16 subbands, however, only require 16 BlBs and another 48 are available for baselineboard stacking. One can thus use 4 BlBs for each of the subbands, quadrupling the number of channels from 128 to 512, which reduces the channel widths from 1MHz to 0.25 MHz over the full 2 GHz frequency range. This method works for any subband bandwidth. Baselineboard stacking can be very useful for spectral line work, as it allows for wider bandwidths for each subband while maintaining frequency resolution. With baselineboard stacking, you can use fewer subbands to cover a set frequency range, thereby minimizing the number of filter edges and resulting sensitivity dropoff. Baselineboard stacking is offered for the August 2012 deadline. For shared risk observations, however, we recommend that observers do not use all 64 baselineboards to allow for observing in the event that one or two baselineboards are not working on any given day. <!-- JUERGEN - might be helpful to have a table that shows the # of channels in a single, dual, and full polarization subband and the # of baselineboard pairs allocated --> === Recirculation === "Recirculation" is a second method for obtaining more spectral channels in a given subband. Recirculation uses the fact that the correlator has more computing capability when the data is averaged in time. The standard correlator dump time is 1s. If this is doubled to 2s, WIDAR can produce twice as many channels as listed in the tables above. 4s would allow 4 times the number of channels. Recirculation is only available for shared risk observations for the August 2012 proposal deadline. == Narrow Subbands with the 3-bit sampler == As the 3-bit samplers are still under commissioning, the August 2012 proposal deadline offers only a single observing setup with 3-bit sampling. This setup includes full coverage over all four baseband pairs (using 64 128MHz subband) 2MHz resolution for full polarization, 1MHz dual, and 0.5MHz for single polarization. == Data Rate Limits == Baselineboard stacking, recirculation, and time resolution can add up to an extremely high data rate in the correlator. The VLA currently supports data rates of up to 20MB/s for regular and more for shared risk observing, see the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary for details]. The OPT instrument configuration calculates data rates based on the spectral line setup and the data rate maxima should not be exceeded for any observational setup. = Tips for Planning, Setup, and Processing of VLA Spectral Line Observations = '''Reminder:''' The following capabilities are offered for regular observing for the August 1 2012 proposal deadline. Consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary] for additional options available through shared risk observing: * Maximum data rates of 20MB/s * Doppler setting * 8-bit samplers providing ** two 1GHz basebands ** up to 16 independently tunable subbands per baseband ** independent number of polarization products in each subband ** independent subband bandwidths ranging from 31.25kHz to 128MHz ** independent number of channels in each subband; maximum number of channels distributed across all subbands and polarization products cannot exceed 16384 channels * 3-bit samplers providing ** four 2GHz basebands with a total of 64 128MHz subbands. 2MHz resolution full polarization, 1MHz dual and 0.5MHz single polarization == Considerations for Planning Subband Bandwidths and Resolution == === Bandwidths required for UV continuum subtraction === When determining the bandwidths needed in your subbands, it is important to observe enough line-free channels on each side of the spectral line to allow for good continuum subtraction. It is possible to interpolate the continuum levels from one subband to another, but it is usually a better solution to derive the continuum level in each subband separately. The number of line free channels ideally equals or exceeds the number of channels that cover the line. When the line free channels are distributed equally on both sides of the spectral line, a low order polynomial (polynomials of order 1 appear to be good models for most cases) usually provides a good fit. Whenever higher order polynomials are needed, e.g. for a continuum source with a significant spectral curvature over the subband(s), the number of line-free channels should be increased. === Channel Widths for High Dynamic Range Imaging: Dealing with the Gibbs Phenomenon === For very sharp spectral or lag features, the spectrum can prominently display a ringing effect known as the Gibbs phenomenon, a sinc function that zig-zags on alternating channels. If this is apparent in the data, smoothing adjacent channels will reduce or even eliminate the effect. Hanning smoothing is the most effective method, which uses a triangular smoothing kernel with the central channel weighed by 0.5 and the two adjacent channels by 0.25. After Hanning smoothing, however, the channels are not independent anymore and one can eliminate every other channel without losing signal to noise. In the pre-upgrade VLA days, the correlator design had a realtively short, truncated lag spectrum, which could result in prominent Gibbs ringing. To avoid this effect, Hanning smoothing was frequently applied online during the observations. With the new WIDAR capabilities of the upgraded VLA, however, ringing is very rare and only observed for extremely strong maser or RFI sources. Consequently, the VLA does not support online Hanning smoothing anymore; if required, Hanning smoothing can be applied during post-processing (e.g. with the CASA task [http://casa.nrao.edu/docs/TaskRef/hanningsmooth-task.html <i>hanningsmooth</i>].). The Gibbs effect can also be reduced by using higher spectral resolution that covers the spectral feature with several channels. In that case, the ringing effects from the individual channels beat against each other, effectively reducing the zig-zag pattern that appears on alternating, neighboring channels when the peak is within a single spectral channel. == Sensitivity/Exposure Time Calculation == [[File:ExposureCalculator.png|300px|thumb|right|The VLA Exposure Calculator]] After planning your general setup and sensitivity needs, the required on-source observing time is best calculated with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA Exposure Calculator]. This JAVA tool (see screenshot) allows one to input the required rms noise and bandwidth limits and outputs the required time on source given a frequency, weather, weighting scheme, and number of polarization products. The input '''Bandwidth''' should correspond to the frequency resolution that is required to perform the science. This may or may not be the width of individual spectral channels. Overheads need to be added according to our [https://science.nrao.edu/facilities/evla/proposing/frequently-asked-questions VLA frequently asked questions webpage]. We refer to the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency low frequency guide] and [http://evlaguides.nrao.edu/index.php?title=Category:HighFrequency high frequency guide] for further advice on how to set up the observations, depending on the receiver band to be used. == The Proposal Submission Tool (PST) == The Proposal Submission Tool (PST, accessible at [http://my.nrao.edu my.nrao.edu] ) is used to submit proposals, including the scientific justification, abstract, and authors, as well as planned target sources, observing session lengths, and correlator setups. In order to ensure that planned correlator setups comply with the capabilities offered for the August 2012 deadline, the Proposal Submission Tool (PST) includes a spectral setup tool. <!-- link to Michaels docs--> [[File:PST-narrow.png|200px|thumb|right|Snapshot of the 8-bit sampler, 2x1GHz PST setup tool]] In the example to the right, 9 subbands were chosen in Ka band with 8-bit sampling. Four of the subbands are in the first and five in the second baseband. The setup features different subband bandwidths, polarization products and uses baselineboard stacking (up to 16 baseline board pairs per subband pair are used for a couple of subbands). A total of 49 baseline boards are used for this configuration. [[File:PST-wide.png|200px|thumb|right|Snapshot of the 3-bit sampler, wideband mode PST setup tool]] Wide-band mode (3-bit sampler, up to 8GHz bandwidth) can be used for spectral line observing as well. The channel width is fixed to a resolution of 2MHz for full, 1MHz for dual and 0.5MHz for single polarization products. No narrow subbands can be chosen. == Setting up a Spectral Observation using the Observation Preparation Tool (OPT) == The Observation Preparation Tool (OPT) is the web-based interface to create scheduling blocks (SBs) for time awarded on the VLA. An SB is the observing program used for a single observing run. This consists of at least a few startup scans, a pointing reference, a bandpass calibration, a flux calibration, gain calibration and target observations. In the OPT, you specify your sources, scan lengths and order, and correlator setups. A full project may consist of several SBs. To access the OPT, go to [http://my.nrao.edu my.nrao.edu] and click on the ''Obs Prep'' tab, followed by [https://e2e.nrao.edu/opt/ ''Login to the Observation Preparation Tool'']. Instructions for using the OPT and for selecting appropriate calibrators are provided in the [http://evlaguides.nrao.edu/index.php?title=Category:OPT-QuickStart OPT QuickStart Guide], and a comprehensive user's manual and up to date information on the OPT are available at [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage]. As such, this guide provides only brief notes on the bandpass and gain calibrators and then focuses on the task of setting up correlator resources. === Bandpass Setup === All observations with the VLA - even those with the goal of observing continuum - require bandpass calibration. When scheduling the bandpass calibration scans within an SB, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged in the data for the rest of the observation. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, <math> S_{cal} \sqrt{t_{cal}}</math> should be greater than <math>S_{obj} \sqrt{t_{obj}}</math>. How many times greater will be determined by one's science goals and the practicalities of the observations, but <math> S_{cal} \sqrt{t_{cal}}</math> should be greater by at least a factor of two. For extremely narrow channels or very weak bandpass calibrators, those typical flux requirements can lead to extremely large integration times. As an alternative one may then chose to reduce the integration time and interpolate in frequency, or to fit a polynomial across all channels in post-processing (''bandtype=BPOLY'' in CASA's [http://casa.nrao.edu/docs/TaskRef/bandpass-task.html bandpass task]. The bandpass calibrator should be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite strong and can often double as bandpass calibrators. At high frequencies (Ku, K, Ka, Q), however, these sources have only moderate flux densities of ~0.5-3 Jy, translating into a potentially noisy bandpass solution. A different, stronger bandpass calibrator should then be observed. <!-- JUERGEN - can you suggest any --> Naturally, all of the above depends on the channel widths and for wide channels the standard flux calibrators may be sufficient even at higher frequencies. In turn, extremely narrow channels may require stronger bandpass calibrators at the low frequency end. Additionally, We have shown that one can transfer the bandpass from a wide subband onto a narrow subband if the wide bandpass frequency range covers that narrow one. This may be good to a level of a few per cent, but we advise to use that option only when absolutely necessary. The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2-4) parts in a thousand over a period of several (~4-8) hours. This should be sufficient for most scientific goals but the bandpasses can be observed several times during an observation for extreme calibration accuracy requirements. A complication can occur when the frequency range of the bandpass is contaminated by other spectral features such as RFI lines or Galactic HI in absorption or emission. There are two basic options to accommodate that situation: * if the feature is narrow, one can simply observe as usual. In post-processing, the narrow feature can be flagged and the frequency gap interpolated by values of nearby channels, or by fitting a polynomial across the bandpass. * for wider contaminating lines, an option is to observe the bandpass at slightly offset frequencies and transfer the bandpass to the target frequency. If a common solution is obtained from two, symmetric offsets at higher and lower frequencies, the solution can be improved. Depending on the choice of offsets and also on the position in the receiver frequency range the error can vary. For 4 MHz offsets close to the HI rest frequency of 1.42GHz, the error is in the per cent range. A guide for CASA is described on [http://casaguides.nrao.edu/index.php?title=Combining_Bandpasses this CASAguides wiki page]. <!-- JUERGEN - some of this should be incorporated into the QuickStart Guide --> === Phase/Complex Gain Calibration === The complex gain (phase/gain) calibration is the same for a spectral line observation as for any other observation. Ideally one should use the same correlator setup for the gain calibrator and the science target. For weak calibrators, however, it is possible to use wider bandwidths for the phase calibrator and then transfer the phases to the source. However, there will be a phase offset between them. The phase offset between the narrow and wide subbands can be determined by observing a strong source (e.g. the bandpass calibrator) and applied in post-processing from the complex gain calibrator to the target sources. A similar method can be used if the complex gain calibrator is observed at a slightly different frequency, e.g. to avoid a contaminating line feature such as Galactic HI. === Correlator Resources Setup === '''For the data taken early 2013 (August 2012 deadline), we will provide a specific new OPT Instrument Configuration layout for the regular and shared risk observing modes. The description below is for the current OSRO and RSRO modes. ''' <!-- JUERGEN - Not sure if you need to explain the current OPT resources, since this guide is designed for the Aug 2012 deadline. Just a thought --> Correctly specifying the WIDAR setup is essential for spectral line observations. In the OPT, click on the ''Instrument Configuration'' tab. The most advanced setting is currently the RSRO setting (''File -> Create New -> RSRO Configuration''). This opens a page for the frequency setup as shown in the figures. Note that you may only have access to the OSRO configuration utility, depending on your history of VLA observations. The OSRO setup is a more restricted version of the RSRO setup. * Enter a name and select the receiver in the top panel. This will adjust the available frequency range described below the drop down menu. The 1dB and 3dB ranges describe the roll-off behavior of the receiver sensitivity at the receiver band edges. If the frequency to be observed is close to the edge frequency of a receiver, one may check if the next higher or lower frequency receiver is more suitable. * ''Baseband Tuning'': Select the position of the basebands. For the ''8bit samplers'', the central frequencies for 2 basebands are to be provided, each with a width of 1GHz. The center frequencies will go into the A0C0 box for frequency 1 and into the B0D0 box for frequency 2. For ''3bit samplers'', 4 basebands are available, each 2 GHz wide. Here, the center frequencies need to placed into the A1C1, A2C2, B1D1, and B2D2 boxes. A1C1 and A2C2 cannot be more than 4GHz apart, and the same restrictions apply to B1D1 and B2D2. Note that additional frequency restrictions may apply and the OPT/ICT will issue a clear warning or error message for those cases. The graphical panel above the input boxes shows the position of the two or four basebands. They are displayed as pairs to accommodate the R and L polarization inputs that may later be converted into single, dual, or full polarization products. * ''Integration Time:'' this defines how data is dumped from the correlator backend into data files. A larger integration time will reduce the data volume. In addition, larger values can be chosen to take advantage of recirculation (only shared risk). On the other hand, time smearing effects, RFI excision, or time resolution may demand smaller integration times. It is important though, to not exceed the maximum data rate of 60MB/s (15MB/s for regular observing) and the integration time parameter is a good way to stay below this threshold for observations that demand large number of spectral channels. * The total data rates are displayed in the ''Configuration Summary''. The same panel also shows the number of baseline boards that are used in the setup. Remember that a maximum of 64 baseline boards are available and make sure that the data rate limits are not exceeded. [[File:OPT-config1.png|200px|thumb|right|OPT - Instrument Configuration: Baseband Settings]] [[File:OPT-config2.png|200px|thumb|right|OPT - Instrument Configuration: Subband Settings]] * '''Subband Setup''': Depending on the baseband setup, the '''Subband Configuration''' panel sports tabs for each baseband. Under each tab one can now select the individual subbands. Up to 64 subbands are available (16 for regular observing): click ''Add subband'' to create a subband setting and select the frequency range from the "Sky Range" drop-down menu. The ''Offset Freq from Center'' shows the placement of the subband with respect to the baseband center. For small bandwidths, the drop-down menu is not available as there are too many choices and the placement needs to be entered by hand in the ''Offset Freq from Center'' box. In the current OSRO/RSRO interface, the subbands are not independently tunable yet (but this feature will be available for the August 2012 deadline) and the subbands will snap on a frequency grid defined by the subband bandwidth. Now select the number of polarization products and the number of channels will be displayed in the ''Spectral Points'' box. The ''comments box'' can be used to describe the setup, e.g. by entering the transitions that should fall in that subband. Those entries are currently not used anywhere outside OPT. The ''delete'' button removes the subband if is no longer required, and ''Bulk Edit'' is used for bulk editing of many subbands (see the OPT guide on this feature). '''Note: if you chose subbands with different bandwidths during OSRO/RSRO, contact NRAO staff as these scripts currently need manual editing after OPT submission.''' * If not all subbands are used, one can use the remaining baseline boards to obtain a higher spectral resolution for those in the Subband Configuration panel. Select a higher number in the ''BL.BPS'' drop-down panel for baselineboard stacking. During commissioning, we recommend to use 2,4,8, etc. BlBs here but in principle any of the options in the drop down menu should work. * Recirculation: only shared risk <!-- JUERGEN - maybe consider giving only a few pointers, and saying that the instructions are in the quick start guide. I'm not sure that you need specific instructions in this guide. --> ==== Using Doppler Setting ==== [[File:Doppler.png|200px|thumb|right|OPT - Doppler Setting]] Doppler setting will calculate the sky frequency of your observation based on the time of the observation, the source velocity, position and line rest frequency. In contrast to Doppler tracking, Doppler setting calculates this once '''for each baseband''' at the start of the observation (execution time of the SB) and the sky frequency will stay fixed for the entire run of the SB. Every subsequent run of the SB will perform a recalculation of the sky frequency. Doppler setting is currently only supported in OSRO mode (and will be available for the August 2012 deadline), RSRO observers need to calculate their own sky frequency (which is usually not too difficult given that the movement of the earth shifts the line by <math>\pm</math>30 km/s over a year, <math>\pm</math>0.5 km/s over a day - a velocity range that can under almost any circumstances be accommodated for by the wide bandwidths of WIDAR). <!-- JUERGEN - most of this is said above. Not sure if you want it to be said twice. --> To use Doppler Setting, first select ''Rest'' in the baseband frequency setup section of the OPT/Instrument Configuration Tool. All velocities will then be calculated against the center baseband frequency which may or may not be the rest frequency of your spectral line. Supply the reference position (a source from a catalog can be chosen, typically this will be your target source) and the velocity with their frames in the ''Doppler Setting'' section in the OPT/ICT. Doppler Setting can only be applied for each baseband and all subbands will shift by the same frequency amount, offset by the the subband tunings. == Post-Processing Guidelines == please see our [http://casaguides.nrao.edu/index.php?title=EVLA_Tutorials extensive VLA turoials] on the [http://casaguides.nrao.edu/ CASAguides wiki] for examples of how to process VLA spectra line data. 1ba495a20e79bf8f2274cd171250209092877109 14 27 1328 904 2012-07-05T16:39:53Z Emomjian 14 /* Introduction, Low Frequency Observing (L, S & C Bands) */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for EVLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst EVLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All EVLA antennas have been equipped with L and C band receivers. As of January 2012, 20 EVLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Sensitivity Observational status summary page]. ===Correlator Setup=== As of September 2011, the Open Shared Risk Observing (OSRO) program offers up to 2 GHz, configured as two independently tunable basebands, each with up to 8 contiguous sub-band pairs of identical bandwidth and channelization. See the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page] for a complete description. You should pick the correlator mode that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the EVLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 0f51f50e48584e938096ebe835354c62881abc46 1329 1328 2012-07-05T17:03:10Z Emomjian 14 /* Proposing for Low Frequency (L,S,C) Observations */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S-band (2-4 GHz), which provide bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, the 3-bit samplers will provide a total frequency coverage of 4 GHz, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) status of the 3-bit samplers, and the their requirement of increased overhead time for setup and calibration during the observations, will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the EVLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] be1ebe88056a89384f6b63462c05f2f3ef354d52 1330 1329 2012-07-05T17:03:42Z Emomjian 14 /* Radio Frequency Interference (RFI) */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S-band (2-4 GHz), which provide bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, the 3-bit samplers will provide a total frequency coverage of 4 GHz, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) status of the 3-bit samplers, and the their requirement of increased overhead time for setup and calibration during the observations, will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The EVLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] d95bae04758c275cd7fd923c4c2fcbb619fbea71 1331 1330 2012-07-05T17:05:09Z Emomjian 14 /* Configuration */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S-band (2-4 GHz), which provide bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, the 3-bit samplers will provide a total frequency coverage of 4 GHz, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) status of the 3-bit samplers, and the their requirement of increased overhead time for setup and calibration during the observations, will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure EVLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 95d6081468c5a93a77b828b950f63c434089056d 1332 1331 2012-07-05T17:07:17Z Emomjian 14 /* Estimating Sensitivities */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S-band (2-4 GHz), which provide bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, the 3-bit samplers will provide a total frequency coverage of 4 GHz, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) status of the 3-bit samplers, and the their requirement of increased overhead time for setup and calibration during the observations, will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the EVLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the EVLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] a8a0ff5d4b6e6d07b3de75012aa368fd833ec377 1333 1332 2012-07-05T17:07:49Z Emomjian 14 /* Bandpass Calibration */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S-band (2-4 GHz), which provide bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, the 3-bit samplers will provide a total frequency coverage of 4 GHz, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) status of the 3-bit samplers, and the their requirement of increased overhead time for setup and calibration during the observations, will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the VLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the VLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry EVLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 2820d63af06f4e0261062ee91bae27505b013505 1334 1333 2012-07-05T17:08:14Z Emomjian 14 /* Polarization Calibration */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S-band (2-4 GHz), which provide bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, the 3-bit samplers will provide a total frequency coverage of 4 GHz, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) status of the 3-bit samplers, and the their requirement of increased overhead time for setup and calibration during the observations, will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the VLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the VLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry VLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html EVLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability the EVLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 9107bf364a1ea88e55f816300daef98065cc1e0d 1335 1334 2012-07-05T17:09:45Z Emomjian 14 /* OPT Observer's Checklist */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S-band (2-4 GHz), which provide bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, the 3-bit samplers will provide a total frequency coverage of 4 GHz, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) status of the 3-bit samplers, and the their requirement of increased overhead time for setup and calibration during the observations, will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the VLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the VLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry VLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html VLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/evla-frequency-bands-and-tunability the VLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 905099b1ce1767d6f417c0b89f4a616df01d3a64 1336 1335 2012-07-05T17:10:05Z Emomjian 14 /* OPT Observer's Checklist */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S-band (2-4 GHz), which provide bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, the 3-bit samplers will provide a total frequency coverage of 4 GHz, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) status of the 3-bit samplers, and the their requirement of increased overhead time for setup and calibration during the observations, will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the VLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the VLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry VLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html VLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/evla-frequency-bands-and-tunability the VLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 1562eb8235badf0596bc4654d98dbb759bab9f27 1337 1336 2012-07-05T17:10:30Z Emomjian 14 /* Estimating Sensitivities */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S,C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S-band (2-4 GHz), which provide bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, the 3-bit samplers will provide a total frequency coverage of 4 GHz, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) status of the 3-bit samplers, and the their requirement of increased overhead time for setup and calibration during the observations, will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the VLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the VLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry VLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html VLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/evla-frequency-bands-and-tunability the VLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] ff594ab79b988cf30ed996570594fecf3b15a07c 1338 1337 2012-07-05T17:38:53Z Emomjian 14 /* Proposing for Low Frequency (L,S,C) Observations */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S, C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Plots of general receiver availability can [https://science.nrao.edu/facilities/evla/earlyscience/osro be found here] . Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S-band (2-4 GHz), which provide bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, the 3-bit samplers will provide a total frequency coverage of 4 GHz, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) status of the 3-bit samplers, and the their requirement of increased overhead time for setup and calibration during the observations, will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the VLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the VLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry VLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html VLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/evla-frequency-bands-and-tunability the VLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 1ea0998a52537caed20284056a22ab6cc24f6072 1339 1338 2012-07-05T19:01:58Z Emomjian 14 /* Observing Frequencies */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S, C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S-band (2-4 GHz), which provide bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, the 3-bit samplers will provide a total frequency coverage of 4 GHz, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) status of the 3-bit samplers, and the their requirement of increased overhead time for setup and calibration during the observations, will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the VLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the VLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry VLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html VLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/evla-frequency-bands-and-tunability the VLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 995130645c055e1d9c45e6a24eeec48897d9bba8 1340 1339 2012-07-05T19:09:45Z Emomjian 14 /* Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S, C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers (enabling up to 8 GHz total bandwidth) at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S (2-4 GHz) bands, which provide bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, the 3-bit samplers will provide a total frequency coverage of 4 GHz, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) status of the 3-bit samplers, and the their requirement of increased overhead time for setup and calibration during the observations, will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the VLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the VLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry VLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html VLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/evla-frequency-bands-and-tunability the VLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] a954f062773d8df1d5dcb63018b0a02e577da0b1 1341 1340 2012-07-06T00:33:52Z Emomjian 14 /* Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S, C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers (enabling up to 8 GHz total bandwidth) at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays, and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S (2-4 GHz) bands, which provide total bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, a total frequency coverage of 4 GHz can be obtained through the 3-bit samplers, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) state of the 3-bit samplers, and their requirement of increased overhead time for setup and calibration during the observations, the use of the 3-bit samplers at C-band will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see Figures 3 and 4 on the [https://science.nrao.edu/facilities/evla/early-science/osro OSRO page]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the VLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the VLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry VLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html VLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/evla-frequency-bands-and-tunability the VLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] cde23115373a824d120d77341ef83a2c536036ab 1342 1341 2012-07-06T00:38:21Z Emomjian 14 /* Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S, C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers (enabling up to 8 GHz total bandwidth) at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays, and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S (2-4 GHz) bands, which provide total bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, a total frequency coverage of 4 GHz can be obtained through the 3-bit samplers, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) state of the 3-bit samplers, and their requirement of increased overhead time for setup and calibration during the observations, the use of the 3-bit samplers at C-band will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see the figures in the following section [https://science.nrao.edu/facilities/evla/observing/spectral-line-observing#section-9] of the Spectral Line Observing guide). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the VLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the VLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry VLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html VLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/evla-frequency-bands-and-tunability the VLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] e9d91d4f6a402198b6e779b9201eddc898372d85 1343 1342 2012-07-06T00:39:20Z Emomjian 14 /* Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S, C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers (enabling up to 8 GHz total bandwidth) at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays, and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S (2-4 GHz) bands, which provide total bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, a total frequency coverage of 4 GHz can be obtained through the 3-bit samplers, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) state of the 3-bit samplers, and their requirement of increased overhead time for setup and calibration during the observations, the use of the 3-bit samplers at C-band will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see the figures in the [following section https://science.nrao.edu/facilities/evla/observing/spectral-line-observing#section-9] of the Spectral Line Observing guide). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the VLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the VLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry VLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html VLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/evla-frequency-bands-and-tunability the VLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 2e4937204765c400b376a562d6efe62184c681de 1344 1343 2012-07-06T00:40:18Z Emomjian 14 /* Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S, C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers (enabling up to 8 GHz total bandwidth) at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays, and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S (2-4 GHz) bands, which provide total bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, a total frequency coverage of 4 GHz can be obtained through the 3-bit samplers, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) state of the 3-bit samplers, and their requirement of increased overhead time for setup and calibration during the observations, the use of the 3-bit samplers at C-band will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see the figures in the following section [https://science.nrao.edu/facilities/evla/observing/spectral-line-observing#section-9 of the Spectral Line Observing guide]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the VLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the VLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry VLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html VLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/evla-frequency-bands-and-tunability the VLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] b9fd3905d1dbd55769cb79d05b946d862f452c50 1345 1344 2012-07-06T00:41:17Z Emomjian 14 /* Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S, C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers (enabling up to 8 GHz total bandwidth) at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays, and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S (2-4 GHz) bands, which provide total bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, a total frequency coverage of 4 GHz can be obtained through the 3-bit samplers, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) state of the 3-bit samplers, and their requirement of increased overhead time for setup and calibration during the observations, the use of the 3-bit samplers at C-band will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see the two spectra in the following section [https://science.nrao.edu/facilities/evla/observing/spectral-line-observing#section-9 of the Spectral Line Observing guide]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The affect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the VLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the VLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry VLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html VLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/evla-frequency-bands-and-tunability the VLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] ec0f0225620857853b8fd4b7dd61e840ae052217 1346 1345 2012-07-06T00:44:33Z Emomjian 14 /* Radio Frequency Interference (RFI) */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S, C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers (enabling up to 8 GHz total bandwidth) at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays, and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S (2-4 GHz) bands, which provide total bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, a total frequency coverage of 4 GHz can be obtained through the 3-bit samplers, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) state of the 3-bit samplers, and their requirement of increased overhead time for setup and calibration during the observations, the use of the 3-bit samplers at C-band will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see the two spectra in the following section [https://science.nrao.edu/facilities/evla/observing/spectral-line-observing#section-9 of the Spectral Line Observing guide]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The effect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the VLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the VLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry VLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html VLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/evla-frequency-bands-and-tunability the VLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 325b764fe1d4431d35322454e314d2d789b4b3f2 1347 1346 2012-07-06T00:49:05Z Emomjian 14 /* Correlator Integration Time */ wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S, C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers (enabling up to 8 GHz total bandwidth) at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays, and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S (2-4 GHz) bands, which provide total bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, a total frequency coverage of 4 GHz can be obtained through the 3-bit samplers, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) state of the 3-bit samplers, and their requirement of increased overhead time for setup and calibration during the observations, the use of the 3-bit samplers at C-band will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see the two spectra in the following section [https://science.nrao.edu/facilities/evla/observing/spectral-line-observing#section-9 of the Spectral Line Observing guide]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The effect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the VLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the VLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry VLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/field-of-view/time-averaging-loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html VLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [https://science.nrao.edu/facilities/evla/docs/manuals/oss/performance-of-the-evla/evla-frequency-bands-and-tunability the VLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 59821c12fdc3d1c6a419ec0f6433da68416b51bf 1348 1347 2012-07-06T14:38:11Z Emomjian 14 wikitext text/x-wiki ==Introduction, Low Frequency Observing (L, S & C Bands)== This document is intended for observers planning Karl G. Jansky Very Large Array (VLA) observations at low frequencies, specifically L (1-2 GHz), S (2-4 GHz), and C bands (4-8 GHz). These three receiver bands share at least some of the same problems and solutions, as compared to higher frequency bands. For an overview, general performance, and some specifics of receiver band (e.g. sensitivity, etc.) of the VLA, consult the [https://science.nrao.edu/facilities/evla/docs/manuals/oss-2013a Observational Status Summary]. <!-- Of the three frequency bands noted above, there will be an emphasis on L-band (1-2 GHz), which is the band most affected by some of the concerns we will address in this document. For example, it is important not to propose for day time observations during times of high solar activity (such as close to year 2013, when solar maximum is expected to occur) for L-band. The high solar activity's impact on the ionosphere can cause your data to become more or less useless.--> Proposers should consider the overall observing strategy in detail at the proposal stage in order to accurately estimate the total time request and expected bandwidth after RFI has been taken into account. Below we begin with a general description of what to consider when proposing for low frequency (L, S, or C band) observations, followed by a more detailed discussion of the observation strategy for these bands, and an Observers Checklist. ==Proposing for Low Frequency (L,S, C) Observations== In this section we will discuss several issues which should be considered when proposing for VLA low frequency (L,S,C) observations. All proposals should be prepared using the [https://science.nrao.edu/facilities/evla/proposing/vlapst VLA Proposal Submission Tool (PST)]. ===Observing Frequencies=== All VLA antennas have been equipped with L and C band receivers. As of January 2012, 20 VLA antennas have S band receivers. The installation of the S-band receivers will be concluded by late Summer 2012. Information on the sensitivity of the receivers over the frequency ranges can be found at the [https://science.nrao.edu/facilities/evla/docs/manuals/oss-2013a Observational status summary page]. ===Choice of Samplers: 8-bit vs. 3 bit, and Correlator Setup=== The general capabilities for the 2013A semester (January 25 - April 22, 2013) are extensive with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers (enabling up to 8 GHz total bandwidth) at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays, and a phased array capability for VLBI. Currently, the 3-bit sampler system is not as reliable as the 8-bit samplers. This is inherent in the sampler hardware. For the L (1-2 GHz) and S (2-4 GHz) bands, which provide total bandwidths of 1 and 2 GHz, respectively, the 8-bit samplers should be the obvious choice. For C-band (4-8 GHz) observations, a total frequency coverage of 4 GHz can be obtained through the 3-bit samplers, while the 8-bit samplers will only deliver up to 2 GHz instantaneous bandwidth. However, with the current (less reliable) state of the 3-bit samplers, and their requirement of increased overhead time for setup and calibration during the observations, the use of the 3-bit samplers at C-band will at best result in a marginal improvement in the RMS noise in spite of doubling the total bandwidth. Therefore, for all three receiver bands covered in this document, L, S, and C, we recommend the use of the 8-bit samplers. For the correlator mode, you should pick the one that best fits your science case. For continuum science, choose the widest bandwidth per subband with coarse spectral resolution. For spectral-line, choose the subband bandwidth and the spectral resolution that better fit the scientific objectives of your project. Issues that should be kept in mind are: * The subbands in a baseband do not overlap. In addition, a few edge channels may need to be flagged because of the higher noise at the subband edges due to filter roll-off. Thus, if you have many spectral lines such that some of them lie near the edges of the subbands, consider tuning the 2nd baseband at a frequency that is offset by 1/2 a subband width with respect to the first baseband. This will give good noise performance across the width of a single baseband (see the two spectra in the following section [https://science.nrao.edu/facilities/evla/observing/spectral-line-observing#section-9 of the Spectral Line Observing guide]). * As described below, each subband must be calibrated independently. Thus if you chose a narrow subband-width for high spectral resolution, you must ensure that you pick flux density and complex gain calibrators that are strong enough for adequate S/N at this subband-bandwidth. See the [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities] section for more information on needed on-calibrator-times for successful calibration. Long slews will be needed to get to suitably strong calibrators. * Note that it will most likely be necessary to Hanning smooth your data in order to get rid of Gibbs ringing (for the theory behind this phenomenon see, for example, http://en.wikipedia.org/wiki/Gibbs_phenomenon). Low frequency observations are prone to strong Radio Frequency Interference (RFI, see below), flagging the RFI will be close to impossible unless you first Hanning smooth your data. This needs to be taken into account when choosing the spectral resolution of your observations, since the effective resolution will be lower than the original, even though the number of channels will stay the same. Note that the frequency resolution (FWHM) of untapered spectra is <math>1.2 \times \Delta\nu</math> (where <math>\Delta\nu</math> is the channel spacing), while the resolution of Hanning tapered spectra is <math>2.0 \times \Delta\nu</math>. * Additionally, for spectral line observations, it is a good idea to pick a spectral resolution that will allow for at least 4-5 channels across your line, given its expected line width. Also here, very narrow and strong spectral features will suffer from 'Gibbs Ringing'. This is easily fixed by Hanning smoothing your data spectrally during data reduction, but you should account for the subsequent decrease in spectral resolution when picking your correlator configuration. ===Radio Frequency Interference (RFI)=== [[Image:3C286_Cband_Darray.jpg|300px|thumb|right|'''Fig. 1.''' RFI in the flux calibrator 3C286 observed in C band, D configuration.]] RFI is a major issue at L, S, and C bands, and should be accounted for properly. Of these three frequency bands, L band (1-2 GHz) is affected the most; up to 40% of this frequency band is contaminated by severe RFI. The effect of the RFI is worst in D-configuration observations. Extended array configurations are generally much less affected by RFI due to phase winding. Moreover, and in general, the RFI is worst at day-time when compared to night-time observations. Also, certain RFI is direction dependent. Therefore, it is important to keep in mind that of the 1 GHz L-band bandwidth, ~600 MHz is the minimum useful bandwidth you will retain after flagging. The S-band (2-4 GHz) is also affected by severe RFI, and out of its 2 GHz bandwidth, ~1500 MHz is the minimum useful bandwidth you will retain after flagging. Continuum observers wishing to use the full 1 GHz range of L-band, or the 2 GHz range of the S-band, should effectively use a narrower bandwidth in the sensitivity calculator ([http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Estimating_Sensitivities Estimating Sensitivities]) while estimating the on-source time and the expected rms noise of their proposed observations, e.g., use 600 MHz and 1500 MHz instead of 1 GHz and 2 GHz, for L- and S-band, respectively. As previously noted, RFI will also affect the spectral resolution of the data – in order to remove the Gibbs ringing of the RFI before flagging, the data will need to be Hanning smoothed. Please take this into account while planning and setting your observations. Detailed information about the RFI at the VLA can be [https://science.nrao.edu/facilities/evla/observing/RFI/index found here]. ===Configuration=== The VLA is reconfigurable. In general, as the baseline length gets longer, the phase stability gets worse. The D-configuration provides the shortest baselines: poorest angular resolution, but highest surface brightness sensitivity. The A-configuration provides the longest baselines: highest angular resolution, but poorest surface brightness sensitivity. See the [https://science.nrao.edu/facilities/evla/proposing/configpropdeadlines configuration schedule] for details for each call for proposals. It is generally important to consider # What angular resolution is required for your science at the desired observing frequency # For resolved sources, how does the desired angular resolution compare to the required surface brightness sensitivity -- particularly important for thermal emission # How much flux will be resolved out by the array configuration that gives the desired angular resolution # Will your source be up during a favorable time of day / month for your observing frequency? The [https://science.nrao.edu/facilities/evla/docs/manuals/oss-2013a/performance-of-the-evla/resolution Observational status summary on resolution] in conjunction with the [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure Exposure Calculator], and the [http://evlaguides.nrao.edu/index.php?title=Monthly_Conditions_at_EVLA Monthly_Conditions_at_EVLA plots] can help answer these questions. Please note that low declination sources risk being subject to antenna shadowing at certain azimuths for the D configuration. <!-- This applies to targets at declination around -17 degrees and lower --> These targets can still be observed, but the lower the declination, the smaller the windows of no-shadowing (one of either side of the north-south arm) will be, which effectively makes the setup of the scheduling blocks harder. ===Estimating Sensitivities=== The [https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure VLA exposure calculator] is an essential tool for estimating the required on-source observing time for the various parts of your observation. Also, see the RFI section above. Below we describe important parameters to estimate various kinds of sensitivities. [[Image:exposure.png|300px|thumb|right|'''Fig. 2.''' Screen capture of the Exposure Calculator for reference.]] '''All Sources''' * ''Configuration'': The configuration choice will not change the point-source rms noise, but it will affect the rms brightness temperature sensitivity which can be important for resolved sources and thermal emission. * For the number of the available antennas at the L, S, and C-bands, check the [https://science.nrao.edu/facilities/evla/docs/manuals/oss-2013a VLA Observational Status Summary]. * ''Elevation'': System temperature increases at low elevations, especially at S-band. Due to ground spillover, S-band observations below 20 degree elevation is not recommended. * ''Average Weather'': This should not be a concern for low frequency (L,S,C band) observations. * Be sure to pick the ''sky frequency'' (not the rest frequency) of your observations. * ''Time on Source'': If you are calculating rms based on a specified time, be sure to enter the appropriate time here. If you are estimating the noise in your science target field over the whole observation, use the total time on source for that field. However, if you are estimating sensitivities for self-calibration be sure to put what you expect the solution interval time to be. At minimum this would be the integration time. '''Continuum Science Target''': * Enter the ''full continuum bandwidth per polarization'' decremented by the approximate bandwidth you expect for any line emission, and in particular RFI (see [http://evlaguides.nrao.edu/index.php?title=Category:LowFrequency#Radio_Frequency_Interference_.28RFI.29 RFI] section above) * Enter ''dual polarization'' '''Spectral Line Science Target''': * Enter the ''channel width'' you plan to use for imaging, i.e. can't be less than the channel width of the selected correlator mode, but could be larger if you plan to average channels during imaging. * Enter ''dual polarization'' (unless you are using a single polarization correlator mode for additional spectral resolution) '''Continuum Calibrator (Absolute Flux, Phase, Amplitude, Pointing):''' * Select ''single polarization'' (each polarization product must be separately calibrated) * Select ''subband width'' of your chosen correlator mode (each subband must be calibrated separately) * For your calibration observations, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. This is typically the scan length at L, S, and C-band, but the specifics of your observations may require a different value. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. '''Spectral Calibrator (Bandpass):''' * Select ''single polarization ''(each polarization product must be separately calibrated) * Select ''channel width'' of your chosen correlator mode (unless you plan to smooth before calibration; though other experimental options exist, typically you calibrate channel by channel) * Again, it is generally a good idea to have a S/N ratio <math> \geq </math> 4 on each baseline of your bandpass calibrator scans. You can use the exposure calculator to determine if the calibrator is appropriate for the selected bandwidth, e.g.: ** Put 2 in the ''Number of Antennas'' field ** Put the appropriate solution interval in the ''Time on Source'' field. ** See if the S/N is <math> \geq </math> 4 within the assumed solution interval. ==Low Frequency Observing Strategy== High solar activity results in emission that can severely affect the data to such a degree as to potentially render your observations useless. Also, such activity causes disturbing ionospheric effects. Therefore, it is imperative to avoid observing at L- and S-bands at day-time (including sunrise and sunset) during times of high solar activity. This is especially important now as we are nearing the max in the solar cycle in 2013. Solar flares with as much as a million Jy at L-band with narrow angular extents are a source of major interference. These flares are equivalent to bright unresolved sources with time-varying flux densities making it very difficult, if not impossible, to remove their effects. As a consequence, the resulting images will be of poor quality and low dynamic range. Some links that can be useful are * Solar activity monitoring: Solar activity and general space weather can be reviewed at the [http://www.swpc.noaa.gov/today.html NOAA site]. The site provides solar activity forecasts and geophysical (geomagnetic field) activity forecasts along with GOES X-ray flux values. *Ionosphere monitoring: [http://iono.jpl.nasa.gov/latest_rti_global.html Global Ionospheric TEC maps]. The quiet Sun also poses problems for L-band observations, particularly in the short configurations (D and C). This would result in degraded image quality if the Sun is too close to the science target to within a few degrees. In the extended configurations (B and A), one can expect the effects of the quiet sun to be reduced. ===Calibrations=== A number of calibrator observations are needed for the subsequent data reductions. These usually include: Flux density calibrator, complex gain (phase and amplitude gain) calibrator, and polarization calibrator if polarization is part of your science objective. You will also need to observe a bandpass calibrator to correct for the delays as well as the relative gains of the spectral channels even if the observations are intended for continuum science. It is recommended that the flux calibrator be observed at least once in an observing run (a scheduling block). The flux calibrator itself, if desired, can be used to correct for the bandpass and antenna based delays, considering that the all the flux calibrators are strong sources at L, S, and C bands. Complex gain calibrators need to be observed before and after each target observation. The time required between phase calibrator observations depends on the array configuration (see below). ====Absolute Flux Calibration==== Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models of these calibrators at L, C, X, Ku, K, and Q bands, are available in both data reduction packages. Flux calibrator models for the S-band will be available in the near future. Meanwhile, observers can utilize the L-band or the C-band models of these calibrators for their S-band data. In general, these calibrators are strong enough to be used for bandpass calibration as well. However, see the following section for the specifics on bandpass calibration, and whether your science goals require observing a stronger bandpass calibrator (i.e., if you are using very narrow channels). Note that bandpass calibration should be carried out even if the observations are intended for continuum science. <!-- Additionally, 0713+438 is a promising unresolved primary calibrator but monitoring data so far is patchy and it is rather weak; 0410+769 is another possibility but monitoring data is patchy for this source as well. 3C84 (0319+415) should be avoided as it is quite variable (but can be used for low polarization leakage calibration). In the future the [http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html polarization monitoring database] will be a useful resource for alternative flux calibrators. --> ====Bandpass Calibration==== For the VLA it is essential to calibrate the spectral response of your chosen correlator mode, even for continuum projects. However, the requirements for bandpass calibration are very dependent on your science goals/type of observation. Below we give some rules of thumb for different kinds of observations. If your project contains a number of aspects (spectral line and continuum), follow the recommendation that is the most limiting. *'''Spectral Line Observations''': Generally speaking, you want to obtain the same rms noise level per channel on your bandpass calibrator as you hope to achieve for your science target. In other words, your bandpass-calibrated science-target spectral-line data will be as noisy as the bandpass calibration applied to it (no matter how sensitive your science target observations originally were). Obtaining sufficient S/N for narrow bandwidth modes can be particularly challenging. For this reason you typically want to observe the strongest source available, preferably in the same part of the sky as your science target to minimize slew times. ** Caveat 1: If you know for certain that you will spectrally average your science target data, you can count on doing this for the bandpass calibrator too. ** Caveat 2: The bandpass response on the VLA antennas are typically smooth over several MHz in bandwidth, suggesting that you can smooth the bandpass calibrator solutions spectrally a bit without averaging over instrumental spectral features that you want to calibrate out. This option is available in both AIPS and CASA. *'''Continuum Observations''': The requirements for bandpass calibration for (pure) continuum observations are significantly less stringent. A signal-to-noise per channel per antenna of about 20 is typically enough to ensure that each subband has adequate bandpass calibration. You may be able to go as low as a S/R of about 10 but this is a function of the science you want to do and should be carefully considered first. Note that instead of using a single observation of a bright bandpass calibrator, it is also possible to use multiple observations (scans within a given scheduling block) of the complex gain calibrator to calibrate the bandpass by combining all its scans. ==== Complex Gain Calibration==== In the choice of the complex gain calibrator, obviously a calibrator close to your target source will decrease slewing times. However, for the L, S, and C bands, we recommend choosing a strong P calibrator further away, rather than a weak nearby S calibrator. See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual] for a description of the calibrator codes P and S. The frequency (in time) of observing a complex gain calibrator in a given scheduling block depends on the configuration of the array. Rough guidelines for the typical time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} ====Polarization Calibration==== There are two components for polarization calibration: # Determining the leakage terms (i.e., the polarization impurity between the R and L polarizations). # Calibrating the absolute polarization angle. There are two approaches to determine the leakage terms: * either observe one or more strong calibrators (> 1 Jy) over a wide range (e.g., > 60 degrees) in parallactic angle and through multiple scans, * or observe a strong unpolarized (typically less than 1% polarized) calibrator source through at least one scan. To calibrate the absolute polarization angle, observe a calibrator with well-known polarization angle. For detailed information on polarization calibration, including the most common calibrators, see the [http://evlaguides.nrao.edu/index.php?title=Category:Polarimetry VLA polarimetry guidelines]. Please note that avoiding day-time observations is even more critical for polarization observations, especially during peak solar activity, due to ionospheric Faraday rotation. ===Correlator Integration Time=== <!--*It is best to use the same correlator integration time for all of your data. This will simplify matters if you decide to further time average later.--> *The correlator integration time should be smaller than the time for significant phase variations, but ideally still yield a self-calibratable S/N on your weakest calibrator. *The shorter the integration time the larger your dataset, so making the integration time unnecessarily short should be avoided. *Amplitude loss due to time averaging (see [https://science.nrao.edu/facilities/evla/docs/manuals/oss-2013a/performance-of-the-evla/field-of-view/time-averaging-loss Observational Status Summary, the section on time-averaging loss]) should be taken into consideration while choosing the correlator integration time. For instance, 1 second in A-configuration, 3 seconds in B- and C-configurations, and 10 seconds in D-configuration, are recommended. ===Overheads and how to split up observations=== The calibrations described above will determine your overheads (except for slew time). However you also need to think about how you will chose to break up your observations. If you have many sources or need good uv-coverage it can be most convenient to observe in one long scheduling block (SB), or a few SBs if your source isn't up long enough for one track. This strategy also minimizes overhead. You can also choose to schedule an observation in smaller SBs, even if the source is up longer than the length of the SB. However, note that this will make the overhead larger. <!--However a long SB also means it is harder to get on the telescope, especially if the project has a lower grade. If so, you can improve your chance of getting on the telescope by splitting your observations into smaller blocks which will be easier to fit into the queue. However, note that this will make the overhead larger.--> The main point is that you need to plan ahead for this. ''Every'' scheduling block will need a complete set of calibrations as described above. Additionally, about 10 minutes is needed at the start of each SB for the setup of the system and moving to your first target. It is important to discuss these choices and the requisite overhead in your technical justification. For example, if your total time request including calibration time is 9 hours and you decide to break this up into 3 x 3-hour SBs, you will need to triple the original time estimate for the bandpass and flux calibrators since each need to be observed in each SB, plus add 30 minutes total for startup (i.e. 10 minutes each SB). Don't forget to take this additional time into account for the total time request at the proposal stage. At these frequency bands, setup time, slew times, and observations of various calibrator sources (flux calibrator, bandpass calibrator, complex gain calibrators, etc...), will add at least 20% to 35% overhead, depending on the specifics of the observations and assuming scheduling blocks of more than 2 hours. If you have many objects at different locations to observe, these times could increased to up to almost 50%. Proposers requesting their observations to be carried out using shorter scheduling blocks should also account for higher overhead percentages. Generally, do not scrimp on overheads, since they are important for adequate calibration. Note that strong calibrators require less on-source time than weak ones. Also, keep in mind that a decrease in time on source of the target will only correspond to a small increase in rms noise (the rms noise is inversely proportional to the square root of time). To achieve a more realistic estimate of the overhead, users can prepare scheduling blocks with various lengths at the time of writing the proposal, using [https://my.nrao.edu/nrao-2.0/secure/Observations.htm the Observation Preparation Tool] <!-- Rough guidelines for maximum time between complex gain calibrator observations are: {| border="1" class="wikitable" style="text-align: center;" |+ |- !Array configuration ! A !! B !! C !! D |- !Time (min) | 10-15 || 20-25 || 30-40 || 50-60 |} --> ==OPT Observer's Checklist== '''Select calibrators''' * Primary flux calibrators: Observe one of the four calibrators (3C286, 3C48, 3C138, or 3C147) to achieve absolute gain calibration. * Complex gain calibrators: from the List of [http://www.vla.nrao.edu/astro/calib/manual/csource.html VLA Calibrators], select compact calibrators that lie as close as possible to your target source (preferably less than 15 degrees for the L, S, and C bands). * The smaller the source-calibrator angular separation, the better. However, in deciding between a nearby but weak 'S' calibrator, and a more distant but strong 'P' calibrator, the stronger 'P' calibrator is usually the better choice. * Bandpass calibrators: Needed for both continuum and spectral line science. Could be the same as the flux calibrator. '''Setup Instrument Configuration:''' * For continuum science: ** Use available NRAO standard setups but consider time averaging; make sure you have short enough integration time in order to reduce time averaging losses. * For Spectral line science: ** Velocity resolution: make sure the spectral resolution is enough to resolve your line (take subsequent Hanning smoothing into account). ** Make sure you have enough line-free bandwidth to measure the continuum flux. ** Correlator integration time: make sure you have short enough integration time in order to reduce time averaging losses, and to offset atmospheric/ionospheric phase winding. '''Check your scheduling block (SB):''' * General: ** Are your sources up? Check the <i>Report</i> of the SB. ** Are the scans sufficiently far away from zenith within the time range you observe? Even though this is not a major concern for L, S, and C band observations, we recommend not to observe within 5 degrees from the zenith. ** Are your source coordinates entered correctly and in the same epoch (J2000)? ** Have you integrated long enough to achieve your desired sensitivity on all sources (i.e., target(s) and calibrators)? * Continuum considerations: ** Are your AC and BD IFs at different (non-overlapping) frequencies? See [https://science.nrao.edu/facilities/evla/docs/manuals/oss-2013a/performance-of-the-evla/evla-frequency-bands-and-tunability the VLA Observations Status Summary] for a description of the IFs AC and BD. ** Is your integration time short enough to reduce time averaging losses? * Spectral line considerations: ** Are the rest frequencies of your lines correct? ** Is the velocity definition correctly specified? ** Is the rest frame correctly specified? ==Post-processing guidelines== There are a number of useful tutorials to study in preparation for the data reductions. See for example: * [http://casaguides.nrao.edu/index.php?title=EVLA_Wide-Band_Wide-Field_Imaging:_G55.7_3.4 Wide-band wide-field L-band D-configuration imaging of supernova remnant G55.7_3.4] * [http://casaguides.nrao.edu/index.php?title=EVLA_6cmWideband_Tutorial_SN2010FZ Calibration and imaging of a single-pointing 6cm (C-band) EVLA wideband continuum dataset on the galaxy NGC2967, D-configuration] * [http://casaguides.nrao.edu/index.php?title=EVLA_Continuum_Tutorial_3C391 Calibration and imaging of a C-band multiple-pointing EVLA continuum dataset on the supernova remnant 3C 391] 8d181ec7dcaca8e07e86841b3da00050ca1d9beb 14 86 1352 1215 2012-09-07T20:48:54Z Akepley 16 /* Putting It All Together: Creating a Scheduling Block */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources? ** Is your scheduling block the correct length (multiple of 30 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 600dc1866dbbc980213b1abf55ee113d0cfd3e0a 1353 1352 2012-09-07T20:51:15Z Akepley 16 /* Checking Your Scheduling Block */ wikitext text/x-wiki ==Introduction== The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 96aabe72b5f158d35e4edb54c5ebbd13dedd78fa 1367 1353 2013-01-22T20:58:54Z Akepley 16 /* Introduction */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still be updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 9426236e0faa4649fa1dd3ba1adda7e2dbb93437 1368 1367 2013-01-22T20:59:19Z Akepley 16 /* Introduction */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 60a2eabc2298cbf678561560fdd360cc5798e83e 1369 1368 2013-01-22T21:00:22Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:59, 12 April 2012 (PDT) 866b8746d6a148a8a360b7dac7e4c00815a4ed28 1370 1369 2013-01-22T21:01:42Z Akepley 16 /* About This Document */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] [[User:Akepley|Amanda Kepley]] ([[User talk:Akepley|talk]]) 13:01, 22 January 2013 (PST) 2d448262c92badde199c9cb461d64753b1c4cc87 1371 1370 2013-01-22T21:02:01Z Akepley 16 /* About This Document */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the EVLA. After reading this guide, you should be able to successfully create a simple EVLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of EVLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total EVLA time allocation may consist of multiple scheduling blocks. Most scheduling blocks are scheduled dynamically based on the conditions at the EVLA. This guide assumes that you understand what data is necessary to successfully calibrate your EVLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary EVLA observational status summary], which gives an overview of the EVLA including calibration, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating EVLA data This quickstart guide does not include the all current observing restrictions for the EVLA. These restrictions are given on the [https://science.nrao.edu/facilities/evla/observing/opt OPT web page] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/evla/observing/opt OPT page] also gives information on the latest features of the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/evla/observing/opt/OPTMANUAL.pdf OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 16e38af055b8ea360efac5e2893ad363aedbd3bf 1372 1371 2013-01-23T16:47:15Z Akepley 16 /* Introduction */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] also gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 70c005fed1b5c67abda90650cfdfd4984472b7e8 1373 1372 2013-01-23T16:49:09Z Akepley 16 /* Introduction */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and ]https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] also gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 5600574eb27518e16f7d4a3ddd350b9fe9e6018c 1374 1373 2013-01-23T16:49:26Z Akepley 16 /* Introduction */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] also gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 68b00a5b21a81206fac259853ac6bf0f4d8aab01 1375 1374 2013-01-23T16:49:56Z Akepley 16 /* Introduction */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the EVLA receivers and correlator for your proposed observations. An individual EVLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your time allocation email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 82b2a70b77ae2c235e947daf83ecf6661718806d 1376 1375 2013-01-23T16:50:30Z Akepley 16 /* Overview of the Scheduling Block Creation Process */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 8db800b6dfcea47a6f9466654aafe84bbb55061e 1377 1376 2013-01-23T16:55:24Z Akepley 16 /* OPT access and layout */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible. This option is useful if you're been lucky enough to get lots of VLA projects, but only need to view the projects you're actively working on. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 7ffdac5c318e1ba36b5e3b638f4ebec69a9b3fed 1378 1377 2013-01-23T16:56:19Z Akepley 16 /* OPT access and layout */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool. This option is useful if you're been lucky enough to get lots of VLA projects, but only need to view the projects you're actively working on. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 4b65b65a9068cb6dd61b12724153121302b062f0 1379 1378 2013-01-23T16:57:24Z Akepley 16 /* OPT access and layout */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool. If you've been lucky enough to have many accepted proposals, this option allows you to view only the projects you're actively working on. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is the legacy ID for your proposal (e.g., AK774). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. To add a source velocity to a newly created source, select File-&gt;Create New-&gt;Source Velocity. This will add a source velocity section to the source information in the editing window. If you want to look at the parameters for or edit the sources in your catalog, click on the edit button to the left of the source name in the catalog as shown in the image below. [[Image:source_edit_button.png]]<br /> Most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing] and for [http://www.vla.nrao.edu/astro/calib/manual/hints.html tips on selecting complex gain calibrators.] Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> You can use a similar procedure as the above to add a source from NED or SIMBAD. Use the "External Search" box to find your source. Once you have found the appropriate source, follow the same procedure as above to copy and paste it into your catalog. ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 402680bfeda43cc1ca4a9dc42b5ad8d218121766 1380 1379 2013-01-23T21:11:59Z Akepley 16 /* Creating a Source List */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool. If you've been lucky enough to have many accepted proposals, this option allows you to view only the projects you're actively working on. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (12B-292). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. By clicking "Add" in the "Source Velocity" section, you can add a source velocity if your source doesn't have one already. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>When you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 4abd777be6fbef1abe0eccaaa0edcf9f77a27bf8 1381 1380 2013-01-23T21:12:51Z Akepley 16 /* Selecting a Complex Gain Calibrator */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool. If you've been lucky enough to have many accepted proposals, this option allows you to view only the projects you're actively working on. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (12B-292). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. By clicking "Add" in the "Source Velocity" section, you can add a source velocity if your source doesn't have one already. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) f4c067b6722bd3ec66976caa00d499850d47f69d 1382 1381 2013-01-23T21:13:24Z Akepley 16 /* Selecting a Flux Calibrator */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool. If you've been lucky enough to have many accepted proposals, this option allows you to view only the projects you're actively working on. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (12B-292). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. By clicking "Add" in the "Source Velocity" section, you can add a source velocity if your source doesn't have one already. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <br/> [[File:Flux_cal.png.png]] Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) b4be07ff980c52b0ee35b0b027d2a6e3ab423ad2 1383 1382 2013-01-23T21:14:03Z Akepley 16 /* Selecting a Flux Calibrator */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool. If you've been lucky enough to have many accepted proposals, this option allows you to view only the projects you're actively working on. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (12B-292). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. By clicking "Add" in the "Source Velocity" section, you can add a source velocity if your source doesn't have one already. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) e8771464a1f96d534636615b8d603960827e882a 1384 1383 2013-01-23T21:38:38Z Akepley 16 /* OPT access and layout */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (12B-292). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. By clicking "Add" in the "Source Velocity" section, you can add a source velocity if your source doesn't have one already. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 50757740f0fc149769a40260946dee486e5e4c94 14 86 1385 1384 2013-01-23T21:41:00Z Akepley 16 /* Adding Your Science Targets */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. By clicking "Add" in the "Source Velocity" section, you can add a source velocity if your source doesn't have one already. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) dd813b146d5f0dfd154de17923b84f755a12d5cb 1386 1385 2013-01-23T21:42:29Z Akepley 16 /* Adding Your Science Targets */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The legacy ID for your proposal is given on the cover page of your proposal or in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) d39b9f354dd41ab6c603dce7150dabac8236541f 1387 1386 2013-01-23T21:43:29Z Akepley 16 /* Adding Your Science Targets */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) af24a70e3fe3f1c3587d0d4ee2399cd9183010a5 1388 1387 2013-01-23T21:44:02Z Akepley 16 /* Adding Your Science Targets */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) c49d757ffc3959e7f84260c998a32232b7193db0 1389 1388 2013-01-23T21:57:43Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the general interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR correlator. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequency range for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual for a more complete description of the correlator modes] and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary] ===Configuring for "Continuum" Observations=== If you observing continuum (or if the default continuum setups adequately resolve your line), there are default continuum configurations available for you to use in the "NRAO Defaults" group. This group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. Once you copy and past the configuration into your personal catalog, you can change the default baseband frequencies, polarization, and bandwidth. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. """As of 1/23/2013, the OPT mode for setting up spectral line modes is still not released. Once it is released, updated documentation will follow.""" ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 647f54913eab3b8f382f4f2dd2640cbafcfd1183 1390 1389 2013-01-23T21:58:44Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the general interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR correlator. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequency range for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual for a more complete description of the correlator modes] and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary] ===Configuring for "Continuum" Observations=== If you observing continuum (or if the default continuum setups adequately resolve your line), there are default continuum configurations available for you to use in the "NRAO Defaults" group. This group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. Once you copy and past the configuration into your personal catalog, you can change the default baseband frequencies, polarization, and bandwidth. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> ===Configuring for Spectral Line Observations=== """As of 1/23/2013, the OPT mode for setting up spectral line modes is still not released. Once it is released, updated documentation will follow.""" If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 35a88be7326b4f9b2c459e9a216c176bdbbb411b 1391 1390 2013-01-23T21:59:06Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the general interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR correlator. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequency range for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual for a more complete description of the correlator modes] and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary] ===Configuring for 3- or 8-bit "Continuum" Observations=== If you observing continuum (or if the default continuum setups adequately resolve your line), there are default continuum configurations available for you to use in the "NRAO Defaults" group. This group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. Once you copy and past the configuration into your personal catalog, you can change the default baseband frequencies, polarization, and bandwidth. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> ===Configuring for Spectral Line Observations=== """As of 1/23/2013, the OPT mode for setting up spectral line modes is still not released. Once it is released, updated documentation will follow.""" If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 45c3ec6a2ef634a483b592886d7aa21463e81347 1392 1391 2013-01-23T22:00:52Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the general interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR correlator. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequency range for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' There are no pure continuum (single channel) WIDAR modes. All data is taken in spectral line mode (i.e., many channels) and thus requires bandpass calibration. See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual for a more complete description of the correlator modes] and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary] ===Configuring for 3- or 8-bit "Continuum" Observations=== If you observing continuum (or if the default continuum setups adequately resolve your line), there are default continuum configurations available for you to use in the "NRAO Defaults" group. This group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. Once you copy and past the configuration into your personal catalog, you can change the default baseband frequencies, polarization, and bandwidth. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> ===Configuring for Spectral Line Observations=== """As of 1/23/2013, the OPT mode for setting up spectral line modes is still not released. Once it is released, updated documentation will follow.""" If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) fa7a0aa88d61ecaaa618478859e5143fa82d4d04 1393 1392 2013-01-23T22:01:19Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the general interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR correlator. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequency range for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' There are no pure continuum (single channel) WIDAR modes. All data is taken in spectral line mode (i.e., many channels) and thus requires bandpass calibration. See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual for a more complete description of the correlator modes] and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary] ===Configuring for 3- or 8-bit "Continuum" Observations=== If you observing continuum (or if the default continuum setups adequately resolve your line), there are default continuum configurations available for you to use in the "NRAO Defaults" group. This group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. Once you copy and past the configuration into your personal catalog, you can change the default baseband frequencies, polarization, and bandwidth. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> ===Configuring for Spectral Line Observations=== """As of 1/23/2013, the OPT mode for setting up spectral line modes is still not released. Once it is released, updated documentation will follow.""" If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 179361bc040c0351c555ec8646ae37ce740017f7 1394 1393 2013-01-23T22:03:11Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the general interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR correlator. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequency range for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz filter boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz filter boundaries.'' There are no pure continuum (single channel) WIDAR modes. All data is taken in spectral line mode (i.e., many channels) and thus requires bandpass calibration. See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual for a more complete description of the correlator modes] and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary] ===Configuring for 3- or 8-bit "Continuum" Observations=== If you observing continuum (or if the default continuum setups adequately resolve your line), there are default continuum configurations available for you to use in the "NRAO Defaults" group. This group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. Once you copy and past the configuration into your personal catalog, you can change the default baseband frequencies, polarization, and bandwidth. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> ===Configuring for Spectral Line Observations=== """As of 1/23/2013, the OPT mode for setting up spectral line modes is still not released. Once it is released, updated documentation will follow.""" If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) c61f2d48083d98ebc1b7ead563575d64b436b591 1395 1394 2013-01-23T22:03:56Z Akepley 16 /* Configuring for 3- or 8-bit "Continuum" Observations */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the general interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR correlator. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequency range for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz filter boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz filter boundaries.'' There are no pure continuum (single channel) WIDAR modes. All data is taken in spectral line mode (i.e., many channels) and thus requires bandpass calibration. See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual for a more complete description of the correlator modes] and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary] ===Configuring for 3- or 8-bit "Continuum" Observations=== If you observing continuum (or if the default continuum setups adequately resolve your line), there are default continuum configurations available for you to use in the "NRAO Defaults" group. This group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. Once you copy and past the configuration into your personal catalog, you can change the default baseband frequencies ("SB Center Freq"), polarization, and bandwidth. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> ===Configuring for Spectral Line Observations=== """As of 1/23/2013, the OPT mode for setting up spectral line modes is still not released. Once it is released, updated documentation will follow.""" If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 5e810fd5795e20368bed24d891430994c9311f19 1396 1395 2013-01-23T22:04:33Z Akepley 16 /* Configuring for Spectral Line Observations */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the general interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR correlator. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequency range for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz filter boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz filter boundaries.'' There are no pure continuum (single channel) WIDAR modes. All data is taken in spectral line mode (i.e., many channels) and thus requires bandpass calibration. See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual for a more complete description of the correlator modes] and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary] ===Configuring for 3- or 8-bit "Continuum" Observations=== If you observing continuum (or if the default continuum setups adequately resolve your line), there are default continuum configurations available for you to use in the "NRAO Defaults" group. This group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. Once you copy and past the configuration into your personal catalog, you can change the default baseband frequencies ("SB Center Freq"), polarization, and bandwidth. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> ===Configuring for Spectral Line Observations=== '''As of 1/23/2013, the OPT mode for setting up spectral line modes is still not released. Once it is released, updated documentation will follow.''' If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Setting up the Instrument Configuration== '''Note (1/22/2012): This section no longer applies to the new version of the OPT. We are updating it.''' Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The EVLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the basic OSRO correlator modes. Although the general interface for configuring the correlator is similar for the RSRO modes, RSRO users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR OSRO mode. In OSRO mode, you have two 1024 MHz wide basebands. In each baseband, you may place up to 8 adjacent sub-bands. The individual sub-bands can have widths ranging from 31.25 kHz to 128 MHz, but all sub-bands must have the same width. ''A critical point to remember is that the edges of each sub-band have reduced sensitivity. Do NOT place lines of interest on sub-band edges.'' You can change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. The catalog name is currently "C/D/Any config 2x 1 GHz Full pol (11B OSRO)". You can use the copy and paste procedure to add one of these configurations to your personal catalog. The Observational Status Summary lists the [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#EVLA_Frequency_Bands_and_Tunability default continuum frequencies]. If you wish to create your own continuum configuration, see below for detailed instructions. If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11]of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. To set up your instrument configuration, * Click on the "Instrument Configuration" link in the navigation strip. * Create a new instrument configuration catalog by going to File-&gt;Create New-&gt;Catalog. * This step will be slightly different depending on either or not you have participated in a RSRO program. ** If you are a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New -&gt;OSRO Configuration. ** If you are not a RSRO participant, you can create a new OSRO configuration by going to File -&gt;Create New Configuration. In the instrument configuration setup screen, make sure you select the appropriate OSRO mode under "Editor" (most likely OSRO, not OSRO-2). Your new correlator setup will be configured for continuum observing at X-band. However, the defaults are not particularly sensible (the basebands overlap and the sub-bands are not the maximum width), so you will want to change them even if you are observing continuum at X-band. We will work our way through the instrument configuration settings starting at the top of instrument configuration editing window. * In the first section of the instrument configuration editing window, ** Give an appropriate name to the configuration. ** <p>Select the appropriate receiver from the menu. You may see many errors at the bottom of the screen when you do this. These errors are usually the result of not having tuned your baseband yet, so you can safely ignore them for now (but not forever).</p> <p>[[Image:correlator_name_band.png]]</p> ** Make sure that you have the appropriate editor selected under "Editor". Most likely you will want OSRO. * It is easier to set the baseband frequencies once the graphical displays in the "Sub-band configuration" section are set up properly. Therefore, we will skip the "Frequencies" section of the configuration editor for now and move to the "Sub-band configuration" section. ** <p>Select the appropriate number of sub-bands, sub-bands widths, and sub-band polarizations. Below the sub-band settings there are three plots. The first shows the location of your basebands. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. The second plot shows the location of the sub-bands in your first baseband. The thin dotted lines indicate the edges of the sub-bands. If you hover your mouse over a particular sub-band, the center frequency of the sub-band will be labeled. The third plot shows the same thing for the second baseband. These plots are helpful for determining whether your correlator configuration is correct.</p><p>[[Image:subband_configuration.png]]</p> * Now comes the difficult part -- tuning the basebands. Go to the "Frequencies" section of the configuration editor. The setup here is slightly different depending on whether you are doing line or continuum observations. ** For continuum observations: *** Select "sky" for the doppler setting mode. *** <p>Enter the desired continuum frequency for each baseband in the "SB Center Freq." box. The sky range for each baseband is shown in "BB Sky Range".</p> <p>[[Image:baseband_frequency.png]]</p> ** For spectral line observations using doppler setting, *** Select "rest" next to the baseband frequency. *** <p>You need to enter the rest frequency of the line for the baseband frequency AND an offset equal to half the width of a sub-band. If you do not enter an offset, you will place the line on the boundary between the 4th and 5th sub-bands (if you have 8 total sub-bands). The doppler setting algorithm uses the baseband center frequency to calculate the sky frequency. In the example below, we are using the doppler setting to tune to the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs. The lines are offset from the center of the baseband (and thus the edges of the sub-bands) by including an offset frequency that is half the width of the sub-bands.</p><p>[[Image:baseband_frequency_line.png]] </p><p>[[Image:subband_configuration_line.png]]</p> *** You need to enter the source information in the "Doppler Setting" section. You should already have this information in the source catalog you created earlier. To access your source catalog, click on the "select source" button and navigate to the appropriate source. Selecting "Configure separately" calculates the doppler correction separately for each baseband. Selecting "Configure together" uses the same doppler correction for both basebands. You will most likely want the "Configure separately" option. [[Image:doppler_setting_source_selection.png]] <br /> ** For spectral line observations entering the sky frequency manually, *** Choose "sky" in the Frequencies section. *** Calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. *** Add an offset equal to 1/2 the width of a single sub-band to the sky frequency. *** Enter the sky frequency plus the offset into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. *** <p>In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> *<p>Finally, you need to enter your integration time in the "Correlator Setup" section in the box labeled "integration time". This time depends on your array configuration and observing frequency. See the [https://science.nrao.edu/facilities/evla/observing/restrictions OSRO restrictions] to select the appropriate value for this parameter. There is additional information on this parameter in [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time-Averaging_Loss "Time Averaging Loss"] and [http://evlaguides.nrao.edu/index.php?title=Observational_Status_Summary#Time_Resolution_and_Data_Rates "Time Resolution and Data Rates"] sections of the Observational Status Summary. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> <p>It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The EVLA produces much more data than the VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible.</p> At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 564a504d9fb9c3c398f7b628f1068a33d23f184b 1397 1396 2013-01-23T22:05:08Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the general interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR correlator. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequency range for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz filter boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz filter boundaries.'' There are no pure continuum (single channel) WIDAR modes. All data is taken in spectral line mode (i.e., many channels) and thus requires bandpass calibration. See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual for a more complete description of the correlator modes] and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary] ===Configuring for 3- or 8-bit "Continuum" Observations=== If you observing continuum (or if the default continuum setups adequately resolve your line), there are default continuum configurations available for you to use in the "NRAO Defaults" group. This group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. Once you copy and past the configuration into your personal catalog, you can change the default baseband frequencies ("SB Center Freq"), polarization, and bandwidth. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> ===Configuring for Spectral Line Observations=== '''As of 1/23/2013, the OPT mode for setting up spectral line modes is still not released. Once it is released, updated documentation will follow.''' If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) fa263d23a87a1c7dcc4a7f7883d9907a0e0bec40 1398 1397 2013-01-23T22:06:09Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the general interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR correlator. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequency range for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz filter boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz filter boundaries.'' There are no pure continuum (single channel) WIDAR modes. All data is taken in spectral line mode (i.e., many channels) and thus requires bandpass calibration. See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual for a more complete description of the correlator modes] and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary] ===Configuring for 3- or 8-bit "Continuum" Observations=== If you observing continuum (or if the default continuum setups adequately resolve your line), there are default continuum configurations available for you to use in the "NRAO Defaults" group. This group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. Once you copy and past the configuration into your personal catalog, you can change the default baseband frequencies ("SB Center Freq"), polarization, and bandwidth. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Extragalactic RRLs happen to be broad enough (100s of km/s) to observe in one of the standard "continuum" modes. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> ===Configuring for Spectral Line Observations=== '''As of 1/23/2013, the OPT mode for setting up spectral line modes is still not released. Once it is released, updated documentation will follow.''' If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] for the latest EVLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated EVLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the EVLA pressure histogram for your observing semester. The EVLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the VLA, a scan loop for the EVLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 7799b637e6f8af3e8c79e439a203e86120665b9b 1399 1398 2013-01-23T22:08:55Z Akepley 16 /* Putting It All Together: Creating a Scheduling Block */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the general interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR correlator. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequency range for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz filter boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz filter boundaries.'' There are no pure continuum (single channel) WIDAR modes. All data is taken in spectral line mode (i.e., many channels) and thus requires bandpass calibration. See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual for a more complete description of the correlator modes] and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary] ===Configuring for 3- or 8-bit "Continuum" Observations=== If you observing continuum (or if the default continuum setups adequately resolve your line), there are default continuum configurations available for you to use in the "NRAO Defaults" group. This group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. Once you copy and past the configuration into your personal catalog, you can change the default baseband frequencies ("SB Center Freq"), polarization, and bandwidth. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Extragalactic RRLs happen to be broad enough (100s of km/s) to observe in one of the standard "continuum" modes. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> ===Configuring for Spectral Line Observations=== '''As of 1/23/2013, the OPT mode for setting up spectral line modes is still not released. Once it is released, updated documentation will follow.''' If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, ADD REQUIREMENTS HERE. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 9151a365cffcda3b2e98531ef385d94b8bc75cc1 1400 1399 2013-01-23T22:10:35Z Akepley 16 /* Configuring for Spectral Line Observations */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. ''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the general interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR correlator. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequency range for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz filter boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz filter boundaries.'' There are no pure continuum (single channel) WIDAR modes. All data is taken in spectral line mode (i.e., many channels) and thus requires bandpass calibration. See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual for a more complete description of the correlator modes] and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary] ===Configuring for 3- or 8-bit "Continuum" Observations=== If you observing continuum (or if the default continuum setups adequately resolve your line), there are default continuum configurations available for you to use in the "NRAO Defaults" group. This group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. Once you copy and past the configuration into your personal catalog, you can change the default baseband frequencies ("SB Center Freq"), polarization, and bandwidth. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Extragalactic RRLs happen to be broad enough (100s of km/s) to observe in one of the standard "continuum" modes. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> ===Configuring for Spectral Line Observations=== '''As of 1/23/2013, the OPT mode for setting up spectral line modes is still not released. Once it is released, updated documentation will follow. In the mean time, consult the full OPT manual.''' If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, ADD REQUIREMENTS HERE. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) abfd8c9cad0f7dd906fd037b47c39aa5ff4141bd 1401 1400 2013-01-23T22:15:06Z Akepley 16 /* Introduction */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. The documentation below outlines a procedure for continuum setups. The new OPT spectral line configuration tool will be released soon.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the general interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult EVLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of the WIDAR correlator. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequency range for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz filter boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz filter boundaries.'' There are no pure continuum (single channel) WIDAR modes. All data is taken in spectral line mode (i.e., many channels) and thus requires bandpass calibration. See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual for a more complete description of the correlator modes] and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary] ===Configuring for 3- or 8-bit "Continuum" Observations=== If you observing continuum (or if the default continuum setups adequately resolve your line), there are default continuum configurations available for you to use in the "NRAO Defaults" group. This group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. Once you copy and past the configuration into your personal catalog, you can change the default baseband frequencies ("SB Center Freq"), polarization, and bandwidth. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF THE TWO SUB-BANDS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Extragalactic RRLs happen to be broad enough (100s of km/s) to observe in one of the standard "continuum" modes. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> ===Configuring for Spectral Line Observations=== '''As of 1/23/2013, the OPT mode for setting up spectral line modes is still not released. Once it is released, updated documentation will follow. In the mean time, consult the full OPT manual.''' If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The EVLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. To calculate the sky frequency of your line manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/). You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. Doppler setting mode is configured using the instrument configuration editor. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user) and that RSRO users currently cannot use doppler setting. ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, ADD REQUIREMENTS HERE. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 96f92d24ce8407f27c2591b4ce090c14db10be21 1416 1401 2013-01-25T13:39:31Z Akepley 16 /* Setting up the Instrument Configuration */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. The documentation below outlines a procedure for continuum setups. The new OPT spectral line configuration tool will be released soon.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. '''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.''' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:spectral_line_correlator_basics.png]] * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change this later on in the spectral line configuration tool. The lines that you've added will show up in the overview plot above the tabs. [[Image:spectral_line_correlator_line_setup.png]] * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. [[Image:spectral_line_correlator_baseband_setup.png]] * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. [[Image:spectral_line_correlator_generate_spws.png]] * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.[[Image:spectral_line_correlator_subband_window.png]] * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line tabs, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. [[Image:spectral_line_correlator_doppler.png]] * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. [[Image:spectral_line_correlator_continuum_windows.png]] * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. [[Image:spectral_line_correlator_validation.png]] ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, ADD REQUIREMENTS HERE. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 8ce53208c985a877dd3e3fa99cf8233a1c21d455 1417 1416 2013-01-25T13:40:53Z Akepley 16 /* Configuring for "Continuum" Observations */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. The documentation below outlines a procedure for continuum setups. The new OPT spectral line configuration tool will be released soon.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time.</p> [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:spectral_line_correlator_basics.png]] * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change this later on in the spectral line configuration tool. The lines that you've added will show up in the overview plot above the tabs. [[Image:spectral_line_correlator_line_setup.png]] * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. [[Image:spectral_line_correlator_baseband_setup.png]] * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. [[Image:spectral_line_correlator_generate_spws.png]] * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.[[Image:spectral_line_correlator_subband_window.png]] * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line tabs, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. [[Image:spectral_line_correlator_doppler.png]] * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. [[Image:spectral_line_correlator_continuum_windows.png]] * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. [[Image:spectral_line_correlator_validation.png]] ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, ADD REQUIREMENTS HERE. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 68096d6e87093eeada8bc52d0100a7bb32165745 1418 1417 2013-01-25T13:41:58Z Akepley 16 /* Configuring for "Continuum" Observations */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. The documentation below outlines a procedure for continuum setups. The new OPT spectral line configuration tool will be released soon.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. [[Image:initial_continuum_changes.png]] You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:spectral_line_correlator_basics.png]] * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change this later on in the spectral line configuration tool. The lines that you've added will show up in the overview plot above the tabs. [[Image:spectral_line_correlator_line_setup.png]] * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. [[Image:spectral_line_correlator_baseband_setup.png]] * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. [[Image:spectral_line_correlator_generate_spws.png]] * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.[[Image:spectral_line_correlator_subband_window.png]] * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line tabs, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. [[Image:spectral_line_correlator_doppler.png]] * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. [[Image:spectral_line_correlator_continuum_windows.png]] * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. [[Image:spectral_line_correlator_validation.png]] ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, ADD REQUIREMENTS HERE. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 94d225b2a060983ae6a24c8e0f50284fdcde3142 1419 1418 2013-01-25T13:42:29Z Akepley 16 /* Configuring for "Continuum" Observations */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. The documentation below outlines a procedure for continuum setups. The new OPT spectral line configuration tool will be released soon.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:spectral_line_correlator_basics.png]] * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change this later on in the spectral line configuration tool. The lines that you've added will show up in the overview plot above the tabs. [[Image:spectral_line_correlator_line_setup.png]] * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. [[Image:spectral_line_correlator_baseband_setup.png]] * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. [[Image:spectral_line_correlator_generate_spws.png]] * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.[[Image:spectral_line_correlator_subband_window.png]] * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line tabs, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. [[Image:spectral_line_correlator_doppler.png]] * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. [[Image:spectral_line_correlator_continuum_windows.png]] * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. [[Image:spectral_line_correlator_validation.png]] ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, ADD REQUIREMENTS HERE. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) bcea3a13e92b335d242a51602514ee7c3c2479b5 1420 1419 2013-01-25T13:43:21Z Akepley 16 /* Configuring for "Continuum" Observations */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. The documentation below outlines a procedure for continuum setups. The new OPT spectral line configuration tool will be released soon.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:spectral_line_correlator_basics.png]] * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change this later on in the spectral line configuration tool. The lines that you've added will show up in the overview plot above the tabs. [[Image:spectral_line_correlator_line_setup.png]] * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. [[Image:spectral_line_correlator_baseband_setup.png]] * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. [[Image:spectral_line_correlator_generate_spws.png]] * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.[[Image:spectral_line_correlator_subband_window.png]] * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line tabs, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. [[Image:spectral_line_correlator_doppler.png]] * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. [[Image:spectral_line_correlator_continuum_windows.png]] * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. [[Image:spectral_line_correlator_validation.png]] ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, ADD REQUIREMENTS HERE. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 0983c880d5fcd340c39b680dad7b6feb91639c84 1421 1420 2013-01-25T13:44:07Z Akepley 16 /* Configuring for "Continuum" Observations */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. The documentation below outlines a procedure for continuum setups. The new OPT spectral line configuration tool will be released soon.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. In the example below, we calculated the sky frequency of the H99alpha (rest frequency = 6.676076GHz) and H107alpha (rest frequency = 5.2937333GHz) RRLs for IC 342 on April 11, 2012 when the source will be at zenith. The sky frequencies are 6.675026131 GHz and 5.292900816 GHz for the H99alpha and H107alpha lines, respectively. We added an offset equal to 1/2 the sub-band width (64 MHz/2 = 32 MHz) to the sky frequency to place the line in a sub-band rather than on a sub-band boundary. The final tuning frequencies are 6.707026131 GHz and 5.324900816 GHz.</p><p>[[Image:dopset_input.png]]</p><p>[[Image:dopset_output.png]]</p><p>[[Image:baseband_frequency_line_dopset.png]]</p><p> [[Image:subband_configuration_line_dopset.png]]</p> Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:spectral_line_correlator_basics.png]] * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change this later on in the spectral line configuration tool. The lines that you've added will show up in the overview plot above the tabs. [[Image:spectral_line_correlator_line_setup.png]] * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. [[Image:spectral_line_correlator_baseband_setup.png]] * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. [[Image:spectral_line_correlator_generate_spws.png]] * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.[[Image:spectral_line_correlator_subband_window.png]] * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line tabs, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. [[Image:spectral_line_correlator_doppler.png]] * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. [[Image:spectral_line_correlator_continuum_windows.png]] * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. [[Image:spectral_line_correlator_validation.png]] ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, ADD REQUIREMENTS HERE. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 1d94a94171f304da27c9e2e1d215213b6264399c 1422 1421 2013-01-25T13:45:28Z Akepley 16 /* Configuring for "Continuum" Observations */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. The documentation below outlines a procedure for continuum setups. The new OPT spectral line configuration tool will be released soon.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:spectral_line_correlator_basics.png]] * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change this later on in the spectral line configuration tool. The lines that you've added will show up in the overview plot above the tabs. [[Image:spectral_line_correlator_line_setup.png]] * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. [[Image:spectral_line_correlator_baseband_setup.png]] * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. [[Image:spectral_line_correlator_generate_spws.png]] * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.[[Image:spectral_line_correlator_subband_window.png]] * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line tabs, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. [[Image:spectral_line_correlator_doppler.png]] * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. [[Image:spectral_line_correlator_continuum_windows.png]] * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. [[Image:spectral_line_correlator_validation.png]] ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, ADD REQUIREMENTS HERE. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) bcb6e7563d5883e1d3cc423035fbc606eff87858 1423 1422 2013-01-25T13:47:53Z Akepley 16 /* Configuring for Spectral Line Observations */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. The documentation below outlines a procedure for continuum setups. The new OPT spectral line configuration tool will be released soon.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change this later on in the spectral line configuration tool. The lines that you've added will show up in the overview plot above the tabs. [[Image:spectral_line_correlator_line_setup.png]] * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. [[Image:spectral_line_correlator_generate_spws.png]] * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.[[Image:spectral_line_correlator_subband_window.png]] * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line tabs, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. [[Image:spectral_line_correlator_doppler.png]] * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. [[Image:spectral_line_correlator_continuum_windows.png]] * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. [[Image:spectral_line_correlator_validation.png]] ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, ADD REQUIREMENTS HERE. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) a45d1b28a9d2ef29dbeb22ae417c8efda456cf2e 1424 1423 2013-01-25T13:49:47Z Akepley 16 /* Configuring for Spectral Line Observations */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. The documentation below outlines a procedure for continuum setups. The new OPT spectral line configuration tool will be released soon.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line tabs, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, ADD REQUIREMENTS HERE. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 5dba480ee45f6204bbf1485be5aafa06f138f3f1 1425 1424 2013-01-25T13:51:56Z Akepley 16 /* Configuring for Spectral Line Observations */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. The documentation below outlines a procedure for continuum setups. The new OPT spectral line configuration tool will be released soon.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, ADD REQUIREMENTS HERE. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) ed56360c35b175c53157e7f18bdc68328c068f79 1426 1425 2013-01-25T13:53:06Z Akepley 16 /* Configuring for Spectral Line Observations */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. The documentation below outlines a procedure for continuum setups. The new OPT spectral line configuration tool will be released soon.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line. The imaginary line can be close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, ADD REQUIREMENTS HERE. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) f88804a71bedf1cfd7af52e3cf98a37b40c5a00e 1427 1426 2013-01-25T14:01:20Z Akepley 16 /* Putting It All Together: Creating a Scheduling Block */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. The documentation below outlines a procedure for continuum setups. The new OPT spectral line configuration tool will be released soon.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line. The imaginary line can be close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, you need scans to set-up the attenuators and the gain slopes and scans to set the requantizer gains. The set-up scans for the attenuators need to occur only once at the beginning of a scheduling block. The requantizer gains scans need to occur after every 8-bit to 3-bit transition and every band change in 3-bit. See [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up/3bit Section 7.1.2] of the Guide to Observing with the VLA document for more details and some examples. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) ce127c1a2a335d5f1b79e54a66f4630e2e4577b4 1428 1427 2013-01-25T14:02:47Z Akepley 16 /* Putting It All Together: Creating a Scheduling Block */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. The documentation below outlines a procedure for continuum setups. The new OPT spectral line configuration tool will be released soon.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line. The imaginary line can be close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, you need scans to set-up the attenuators and the gain slopes and scans to set the requantizer gains. The set-up scans for the attenuators need to occur only once at the beginning of a scheduling block. The requantizer gains scans need to occur after every 8-bit to 3-bit transition and every band change in 3-bit. See [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up/3bit Section 7.1.2] of the Guide to Observing with the VLA document for more details and some examples. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. A table to the right of the plot shows the rise/set times for different elevations for your source. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 119397c6cbdbe65e26fc565733b16243f6158f08 1429 1428 2013-01-25T14:04:46Z Akepley 16 /* Putting It All Together: Creating a Scheduling Block */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. The documentation below outlines a procedure for continuum setups. The new OPT spectral line configuration tool will be released soon.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line. The imaginary line can be close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, you need scans to set-up the attenuators and the gain slopes and scans to set the requantizer gains. The set-up scans for the attenuators need to occur only once at the beginning of a scheduling block. The requantizer gains scans need to occur after every 8-bit to 3-bit transition and every band change in 3-bit. See [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up/3bit Section 7.1.2] of the Guide to Observing with the VLA document for more details and some examples. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. A table to the right of the plot shows the rise/set times for different elevations for your source. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). You can also ask to avoid sunrise and sunset, which is important for high frequency observations. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==Scheduling Block Examples== * [https://science.nrao.edu/facilities/evla/observing/opt/lowfreq Low Frequency] (C band)<br /> * [https://science.nrao.edu/facilities/evla/observing/opt/highfreq High Frequency] (Ka band) ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 8671d39d9cacccdc06850faca510bf09dcfe865c 1430 1429 2013-01-25T14:06:27Z Akepley 16 /* Scheduling Block Examples */ wikitext text/x-wiki ==Introduction== '''Note (1/22/2012): This guide is still being updated for the new version of the OPT. The source catalog and scheduling block creation process will still be similar to what's described below. The correlator configuration creation, however, has changed dramatically. The documentation below outlines a procedure for continuum setups. The new OPT spectral line configuration tool will be released soon.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line. The imaginary line can be close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, you need scans to set-up the attenuators and the gain slopes and scans to set the requantizer gains. The set-up scans for the attenuators need to occur only once at the beginning of a scheduling block. The requantizer gains scans need to occur after every 8-bit to 3-bit transition and every band change in 3-bit. See [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up/3bit Section 7.1.2] of the Guide to Observing with the VLA document for more details and some examples. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. A table to the right of the plot shows the rise/set times for different elevations for your source. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). You can also ask to avoid sunrise and sunset, which is important for high frequency observations. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 885670daa7b61ecfe545e733def8697897e5260e 1431 1430 2013-01-25T14:08:03Z Akepley 16 /* Introduction */ wikitext text/x-wiki ==Introduction== '''Note (1/25/2012): This guide has been updated for the lastest version of the OPT. Please contact Amanda Kepley (akepley@nrao.edu) for comments/questions suggestions about this guide and the NRAO helpdesk for any questions about the OPT.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line. The imaginary line can be close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, you need scans to set-up the attenuators and the gain slopes and scans to set the requantizer gains. The set-up scans for the attenuators need to occur only once at the beginning of a scheduling block. The requantizer gains scans need to occur after every 8-bit to 3-bit transition and every band change in 3-bit. See [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up/3bit Section 7.1.2] of the Guide to Observing with the VLA document for more details and some examples. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. A table to the right of the plot shows the rise/set times for different elevations for your source. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). You can also ask to avoid sunrise and sunset, which is important for high frequency observations. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the EVLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) d6442ed25f5a044d5b6ce434cd7537c39b8f6c6c 1432 1431 2013-01-25T14:12:19Z Akepley 16 /* Submitting Your Scheduling Block */ wikitext text/x-wiki ==Introduction== '''Note (1/25/2012): This guide has been updated for the lastest version of the OPT. Please contact Amanda Kepley (akepley@nrao.edu) for comments/questions suggestions about this guide and the NRAO helpdesk for any questions about the OPT.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line. The imaginary line can be close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, you need scans to set-up the attenuators and the gain slopes and scans to set the requantizer gains. The set-up scans for the attenuators need to occur only once at the beginning of a scheduling block. The requantizer gains scans need to occur after every 8-bit to 3-bit transition and every band change in 3-bit. See [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up/3bit Section 7.1.2] of the Guide to Observing with the VLA document for more details and some examples. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. A table to the right of the plot shows the rise/set times for different elevations for your source. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). You can also ask to avoid sunrise and sunset, which is important for high frequency observations. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the EVLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the EVLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the VLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) a8629416eb10157a36241db0b77791ffe6bc1d72 1433 1432 2013-01-25T14:12:50Z Akepley 16 /* Checking Your Scheduling Block */ wikitext text/x-wiki ==Introduction== '''Note (1/25/2012): This guide has been updated for the lastest version of the OPT. Please contact Amanda Kepley (akepley@nrao.edu) for comments/questions suggestions about this guide and the NRAO helpdesk for any questions about the OPT.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line. The imaginary line can be close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, you need scans to set-up the attenuators and the gain slopes and scans to set the requantizer gains. The set-up scans for the attenuators need to occur only once at the beginning of a scheduling block. The requantizer gains scans need to occur after every 8-bit to 3-bit transition and every band change in 3-bit. See [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up/3bit Section 7.1.2] of the Guide to Observing with the VLA document for more details and some examples. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. A table to the right of the plot shows the rise/set times for different elevations for your source. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). You can also ask to avoid sunrise and sunset, which is important for high frequency observations. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the VLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the VLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the VLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 801cc8d97e9c6625ec67b10d1ba659a09eac801f 6 131 1402 2013-01-24T22:44:53Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 132 1403 2013-01-24T22:45:27Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 133 1404 2013-01-24T22:48:26Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 134 1405 2013-01-24T22:49:52Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 1408 1405 2013-01-24T22:52:32Z Akepley 16 Akepley uploaded a new version of &quot;[[File:Spectral line correlator baseband setup.png]]&quot; wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 135 1406 2013-01-24T22:50:56Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 136 1407 2013-01-24T22:51:39Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 137 1409 2013-01-24T22:53:07Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 138 1410 2013-01-24T22:53:29Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 139 1411 2013-01-24T22:53:47Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 140 1412 2013-01-24T22:54:04Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 141 1413 2013-01-25T13:22:30Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 142 1414 2013-01-25T13:22:47Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 6 143 1415 2013-01-25T13:23:04Z Akepley 16 wikitext text/x-wiki da39a3ee5e6b4b0d3255bfef95601890afd80709 0 144 1437 2013-07-17T16:18:35Z Gvanmoor 7 Gvanmoor moved page [[Observational Status Summary - Current]] to [[Observational Status Summary May 20, 2010. Warning: Content Obsolete]] wikitext text/x-wiki #REDIRECT [[Observational Status Summary May 20, 2010. Warning: Content Obsolete]] 75433d014c9d4555409a8b3ba349c8657462cba4 0 125 1438 1223 2013-07-17T19:54:43Z Gvanmoor 7 wikitext text/x-wiki '''REPLACED ON 23 MAY 2012. For an overview of current and past versions see https://science.nrao.edu/facilities/vla/oss/oss''' '''WE RECOMMEND YOU UPDATE YOUR BOOKMARKS ACCORDINGLY''' ---- '''The EVLA Observational Status Summary''' ''Version date: November 29, 2011'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget at the end of 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || align='center'| 10 μJy || align='center'| 1 μJy || align='center'| 10 |- | Maximum BW in each polarization || align='center'| 0.1 GHz || align='center'| 8 GHz || align='center'| 80 |- | Number of frequency channels at max. BW || align='center'| 16 || align='center'| 16,384 || align='center'| 1024 |- | Maximum number of freq. channels || align='center'| 512 || align='center'| 4,194,304 || align='center'| 8192 |- | Coarsest frequency resolution || align='center'| 50 MHz || align='center'| 2 MHz || align='center'| 25 |- | Finest frequency resolution || align='center'| 381 Hz || align='center'| 0.12 Hz || align='center'| 3180 |- | Number of full-polarization sub-correlators || align='center'| 2 || align='center'| 64 || align='center'| 32 |- | Log (Frequency Coverage over 1–50 GHz) || align='center'| 22% || align='center'| 100% || align='center'| 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date !Status or Actual Date |- | Installation of EVLA correlator subset for early science || align='center'| 2010 Q1 || align='center' | 2012 Q1 |- | Shared Risk Observing begins || align='center'| 2010 Q1 || align='center' | 2012 Q2 |- | Last antenna retrofitted || align='center'| 2010 Q2 || align='center' | 2010 Q2 |- | Full EVLA correlator installation || align='center'| 2011 Q2 || align='center' | 2011 Q2 |- | Last receiver installed || align='center'| 2012 Q4 || align='center' | on schedule |- |} == VLA to EVLA Transition == The year 2010 was extremely exciting for the EVLA. The correlator that was the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from '''A'''→'''B'''→'''C'''→'''D'''→'''A''' to '''D'''→'''C'''→'''B'''→'''A'''→'''D''', in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. The last VLA antenna was retrofitted to EVLA specifications in May 2010. During 2011 the WIDAR correlator was put into full observing mode with commissioning and Resident Shared Risk Observing starting in early 2011. By the end of 2011, up to 2 GHz of bandwidth was provided to the the general public along with 2 tunable bands, each with 8 spectral windows (64 to 256 channels) at S, Ku and X bands. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it has not yet been commissioned. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted '''D''', '''C''', '''B''', and '''A''' respectively. In addition, there are 3 “hybrid” configurations labelled '''DnC''', '''CnB''', and '''BnA''', in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle was modified during 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule for 2011, 2012 and 2013 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Approximate EVLA Configuration Schedule for 2011-2012''' ! Year ! Jan-Apr ! May-Aug ! Sep-Dec |- | align='center'| 2011 || align='center'| '''B''' || align='center'| '''A''' || align='center'| '''D''' |- | align='center'| 2012 || align='center'| '''C''' || align='center'| '''B''' || align='center'| '''A''' |- | align='center'| 2013 || align='center'| '''D''' || align='center'| '''C''' || align='center'| '''B''' |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 2012 (a full '''D'''→'''A''' configuration cycle) are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science is provided by two programs for outside users and one for EVLA commissioning staff. All early science programs are peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period thus involves an element of risk associated with the large stepwise increases in throughput bandwidth that are offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program provides early science capabilities to the general user community. These capabilities initially provided a maximum 256 MHz bandwidth that increased to 2 GHz for the '''D''' configuration in mid-2011 and will increase further to 8 GHz at the end of 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All EVLA antennas are outfitted with either EVLA or “interim” L, EVLA C, VLA X, EVLA K, EVLA Ka, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of July 2012, 25 of the EVLA antennas will be outfitted with S-band receivers, 23 EVLA antennas will have Ku-band receivers, and 22 will have new EVLA-style X-band receivers. Figure 1 shows the expected installation rate of final EVLA receiver systems for the rest of the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers, expected to be commissioned by the end of 2012. Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:RcvrAvailDec12.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of final EVLA receivers during the calendar year 2012 until the end of the EVLA Construction Project on December 31, 2012. Interim receivers with reduced frequency coverage or polarization purity are available at some bands in addition to those plotted (see Table 4).]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that are available in September 2011 and January 2012, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, Jan 2012 ! colspan="2" | Receiver availability, Jul 2012 |- ! ! GHz ! Final EVLA ! EVLA+VLA/interim ! Final EVLA ! EVLA+VLA/interim |- | 400 cm (4-band) || align='center'| 0.062–0.078 || align='center'| - || align='center'| - || align='center'|- || align='center'|- |- | 20 cm (L) || align='center'| 1.0–2.0 || align='center'| 19 || align='center'| 27 || align='center'|24 || align='center'|27 |- | 13 cm (S) || align='center'| 2.0–4.0 || align='center'| 21 || align='center'| 21 || align='center'|25 || align='center'|25 |- | 6 cm (C) || align='center'| 4.0–8.0 || align='center'| 27 || align='center'| 27 || align='center'|27 || align='center'|27 |- | 3 cm (X) || align='center'| 8.0–12.0 || align='center'| 16 || align='center'| 27 || align='center'|22 || align='center'|27 |- | 2 cm (Ku) || align='center'| 12.0–18.0 || align='center'| 18 || align='center'| 18 || align='center'|23 || align='center'|23 |- | 1.3 cm (K) || align='center'| 18.0–26.5 || align='center'| 27 || align='center'| 27 || align='center'|27 || align='center'|27 |- | 1 cm (Ka) || align='center'| 26.5–40.0 || align='center'| 27 || align='center'| 27 || align='center'|27 || align='center'|27 |- | 0.7 cm (Q) || align='center'| 40.0–50.0 || align='center'| 27 || align='center'| 27 || align='center'|27 || align='center'|27 |- |} :Note: The "EVLA+VLA/interim" columns give the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers (see Figure 1). == Open Shared Risk Observing (OSRO) == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program has extended this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program obtain good quality data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing (RSRO) == The WIDAR correlator and the EVLA provides a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program is now expected to run through the end of 2012, with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. Full operations of the EVLA will begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program for the upcoming '''D'''→'''A''' configuration cycle, Sep 2011 through Dec 2012. Users interested in participating in the RSRO program should refer to the web page at https://science.nrao.edu/facilities/evla/early-science/rsro for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24 || 80 || 260 || 850 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka) || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800 || 2200 || 20000 || 20000 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka) || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations ('''DnC''', '''CnB''', '''BnA''') should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., '''DnB''', or '''CnA''') is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. For the '''C''' configuration an antenna from the middle of the north arm is moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Note that although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration. ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are insufficient EVLA 8–12 GHz receivers available yet to determine system performance across the band. A project with the goal of doubling the longest baseline available in the '''A''' configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the ''SEFD'' at some fiducial EVLA frequencies. [[File:SEFD.png|none|frame|Figure 2: ''SEFD'' for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the Ku, K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: ''SEFD''s and '''D'''-Configuration Confusion Limits''' !Frequency !''SEFD'' !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || align='center'| 420 || align='center'| 89 |- | 3.0 GHz (S) || align='center'| 370 || align='center'| 14 |- | 6.0 GHz (C) || align='center'| 310 || align='center'| 2.3 |- | 10.0 GHz (X) || align='center'| 250 || align='center'| negligible |- | 15 GHz (Ku) || align='center'| 350 || align='center'| negligible |- | 22 GHz (K) || align='center'| 560 || align='center'| negligible |- | 33 GHz (Ka) || align='center'| 730 || align='center'| negligible |- | 45 GHz (Q) || align='center'| 1400 || align='center'| negligible |} :Note: ''SEFD''s at Ku, K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in '''D''' configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temperature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted above assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. '''At source elevations greater than 80 degrees (zenith angle < 10 degrees), source tracking becomes difficult; it is recommended to avoid such source elevations during the observation preparation setup'''. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in Right Ascension (RA). The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where ''T''_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and ''S'' in mJy per beam, the constant ''F'' depends only upon array configuration and has the approximate value ''F'' = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for ''S''. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of ''F'' calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where ''D'' is the distance to the galaxy in Mpc, and ''S∆V'' is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different “baseband” frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” Currently, a maximum of 1.024 GHz can be correlated for each IF pair (see [[#Correlator Configurations|Correlator Configurations]]), for a total maximum bandwidth of approximately 2 GHz. The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. At X-band a number of the antennas will continue to have the old narrow-band VLA X-band receivers until their retrofit is complete at the end of the EVLA construction project. As of December 2010 there is not a sufficient number of EVLA-style X-band receivers in the array to evaluate either the system performance or the radio frequency interference environment throughout the EVLA X-band tuning range of 8-12 GHz. A total bandwidth of 800 MHz equivalent to that of the VLA receivers (8.0-8.8 GHz) should be assumed for the purposes of sensitivity calculations at X-band for the present. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || align='center'| 1.0–2.0 || align='center'| 1.25¹ || align='center'| 1.75¹ |- | 13 cm (S) || align='center'| 2.0–4.0 || align='center'| 2.5 || align='center'| 3.5 |- | 6 cm (C) || align='center'| 4.0–8.0 || align='center'| 5.0 || align='center'| 6.0 |- | 3 cm (X) || align='center'| 8.0–12.0 || align='center'| 8.5² || align='center'| 9.5² |- | 2 cm (Ku) || align='center'| 12.0–18.0 || align='center'| 13.5 || align='center'| 14.5 |- | 1.3 cm (K) || align='center'| 18.0–26.5 || align='center'| 20.7 || align='center'| 21.7 |- | 1 cm (Ka) || align='center'| 26.5–40.0 || align='center'| 31.5 || align='center'| 32.5 |- | 0.7 cm (Q) || align='center'| 40.0–50.0 || align='center'| 40.5 || align='center'| 41.5 |} :Notes: :: 1. This default frequency set-up for L-band comprises two 512 MHz basebands (each with 8 subbands of 64 MHz) to cover the entire 1-2 GHz of the L-band receiver. :: 2. Many of the antennas continue to have the old narrow-band VLA receivers (see Figure 1), for which a total bandwidth of 800 MHz should be assumed (8.0-8.8 GHz). The RFI environment of the default tunings has not yet been evaluated. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 in [[#Documentation|Documentation]] for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! (∆ν/ν₀)*(θ₀/θ_HPBW) ! Peak ! Width |- | 0.0 || align='center'| 1.00 || align='center'| 1.00 |- | 0.50 || align='center'| 0.95 || align='center'| 1.05 |- | 0.75 || align='center'| 0.90 || align='center'| 1.11 |- | 1.0 || align='center'| 0.80 || align='center'| 1.25 |- | 2.0 || align='center'| 0.50 || align='center'| 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 (in [[#Documentation|Documentation]]) considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || align='center'| 2.1 || align='center'| 4.8 || align='center'| 6.7 |- | '''B''' || align='center'| 6.8 || align='center'| 15.0 || align='center'| 21.0 |- | '''C''' || align='center'| 21.0 || align='center'| 48.0 || align='center'| 67.0 |- | '''D''' || align='center'| 68.0 || align='center'| 150.0 || align='center'| 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 (http://www.aoc.nrao.edu/evla/geninfo/memoseries/evlamemo67.pdf) for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second, for '''A''' configuration. For '''B''', '''C''', and '''D''' configurations the minimum integration time is 3 seconds, unless a shorter integration time is explicitly requested and justified. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 14.4 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 51.8 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 1.2 TB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool (https://science.nrao.edu/facilities/evla/data-archive/evla) will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on disk drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the '''D''' configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum at L-band. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots and tables of known RFI are available online, at http://science.nrao.edu/evla/observing/RFI/; plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the '''D''' configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in '''A''' configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in '''A''' configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List (http://www.vla.nrao.edu/astro/calib/manual/) should be used. The positions of these sources are taken from lists published by the United States Naval Observatory (USNO). == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1 in [[#Documentation|Documentation]]. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth smearing and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its '''D''' configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP (Wilkinson Microwave Anisotropy Probe) flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby calibrator with an "S" code in the calibrator database, and a more distant calibrator with a "P" code, the nearby calibrator is usually the better choice (see http://www.vla.nrao.edu/astro/calib/manual/key.html for a description of calibrator codes). :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173 (http://www.vla.nrao.edu/memos/sci/). These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API may be found at http://www.vla.nrao.edu/astro/guides/api/. Plots of current/historical data can be found at: https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi Characteristic seasonal averages are represented below: {| border="1" align="center" |+ '''Table: Seasonal API/wind values at the EVLA''' !Month !API (night) [deg] !API (median) [deg] !API (day) [deg] !Wind (night) [m/s] !Wind (median) [m/s] !Wind (day) [m/s] |- | [[Media:APIwind_January.png| January]] || 2.3 || 2.8 || 3.6 || 1.6 || 1.9 || 2.3 |- | [[Media:APIwind_February.png| February]] || 2.9 || 3.4 || 4.5 || 4.0 || 4.3 || 4.5 |- | [[Media:APIwind_March.png| March]] || 2.8 || 3.7 || 5.5 || 3.4 || 3.9 || 4.7 |- | [[Media:APIwind_April.png| April]] || 3.3 || 4.5 || 6.2 || 5.3 || 5.5 || 5.8 |- | [[Media:APIwind_May.png | May]] || 2.9 || 4.6 || 6.7 || 2.6 || 3.2 || 3.7 |- | [[Media:APIwind_June.png| June]] || 3.8 || 5.5 || 7.4 || 2.5 || 3.9 || 6.3 |- | [[Media:APIwind_July.png| July]] || 6.2 || 8.3 || 10.5 || 2.9 || 2.9 || 3.0 |- | [[Media:APIwind_August.png| August]] || 5.4 || 7.1 || 11.3 || 1.7 || 2.3 || 3.0 |- | [[Media:APIwind_September.png| September]] || 5.2 || 6.6 || 8.8 || 2.3 || 3.0 || 3.6 |- | [[Media:APIwind_October.png| October]] || 4.2 || 5.3 || 7.4 || 2.3 || 2.9 || 3.7 |- | [[Media:APIwind_November.png| November]] || 2.6 || 3.0 || 4.0 || 1.2 || 2.5 || 1.6 |- | [[Media:APIwind_December.png| December]] || 2.8 || 3.2 || 4.1 || 1.2 || 1.6 || 2.7 |- |} Click on the Month links above to see plots of phase and wind speed versus time. :Note: day indicates sunrise to sunset values; night indicates sunset to sunrise values. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/<math>{\rm{m}^2}</math>, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. For more details see: https://science.nrao.edu/facilities/evla/early-science/polarimetry == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. For the Open Shared Risk Observing (OSRO) program available to the community during the period Sep 2011 through Jan 2013 are offering two independently tunable basebands, where each baseband has up to eight sub-bands. Possible sub-band widths are 128 MHz, 64 MHz, 32 MHz, all the way down in factors of 2 to 0.03125 MHz. All sub-bands must have the same bandwidth and channelization in both basebands, and be contiguous in frequency within each baseband. We will offer three different OSRO modes: full polarization, dual polarization, and single polarization, with 64, 128, and 256 channels per sub-band, respectively. There is always the possibility during offline processing to smooth in frequency to reduce dataset sizes or to improve spectral response. Starting with the '''D'''-configuration in September 2011, we have been providing options for configuring WIDAR for OSRO in the following three ways: :1. “OSRO Full Polarization”: Four polarization products. This configuration offers 4 polarization products for each sub-band, each of which has 128 MHz bandwidth with 64 channels. It is possible to decrease the subband bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in the following table: {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO Full Polarization)''' ! Sub-band BW (MHz) ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 64 || 1000 || 300 || 19,200 |- | 32 || 64 || 500 || 150 || 9,600 |- | 16 || 64 || 250 || 75 || 4,800 |- | 8 || 64 || 125 || 37.5 || 2,400 |- | 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 64 || 31.25 || 9.4 || 600 |- | 1 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO Dual Polarization”: Two polarization products. This configuration offers 2 polarization products for each sub-band, each of which has 128 MHz bandwidth with 128 channels. It is possible to decrease the sub-band bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in the following table. {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for dual polarization (OSRO Dual Polarization)''' ! Sub-band BW (MHz) ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 128 || 1000 || 300/ν (GHz) || 38,400/ν (GHz) |- | 64 || 128 || 500 || 150 || 19,200 |- | 32 || 128 || 250 || 75 || 9,600 |- | 16 || 128 || 125 || 37.5 || 4,800 |- | 8 || 128 || 62.5 || 19 || 2,400 |- | 4 || 128 || 31.25 || 9.4 || 1,200 |- | 2 || 128 || 15.625 || 4.7 || 600 |- | 1 || 128 || 7.813 || 2.3 || 300 |- | 0.5 || 128 || 3.906 || 1.2 || 150 |- | 0.25 || 128 || 1.953 || 0.59 || 75 |- | 0.125 || 128 || 0.977 || 0.29 || 37.5 |- | 0.0625 || 128 || 0.488 || 0.15 || 18.75 |- | 0.03125 || 128 || 0.244 || 0.073 || 9.375 |} :3: "OSRO Single Polarization": One polarization product (new for OSRO observing). It offers 1 polarization product for each sub-band, each of which has 128 MHz bandwidth with 256 channels. It is possible to decrease the sub-band bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in the following table. {| border="1" align="center" |+ '''Table 14: Correlator capabilities per sub-band for single polarization (OSRO Single Polarization)''' ! Sub-band BW (MHz) ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 256 || 250 || 75 || 19,200 |- | 32 || 256 || 125 || 37.5 || 9,600 |- | 16 || 256 || 62.5 || 19 || 4,800 |- | 8 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities are being provided with integration times no shorter than 1 second in '''A''' configuration (3 seconds in '''B'''/'''C'''/'''D''' configurations), and Doppler setting will be available with these correlator configurations. If it is likely that the data will need to be resampled spectrally in order to Doppler track to a line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna (e.g., Y1) modes, have not yet been commissioned and are not yet available to the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing "Faint Images of the Radio Sky at Twenty-centimeters survey(FIRST, http://www.cv.nrao.edu/first/) or the Co-Ordinated Radio 'N' Infrared Survey for High-mass star formation (CORNISH, http://www.ast.leeds.ac.uk/Cornish/public/index.php) ('''B''' configuration), or the NRA VLA Sky Survey (NVSS, http://www.cv.nrao.edu/nvss/) ('''D''' configuration, all-sky) surveys. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 ([[#Documentation|Documentation]]) for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to obtain EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may find that such dissertations comprise pieces of several short proposals, which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being impaired by an adverse review of one proposal when the full scope of the project is not seen. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. Starting in 2011 time on the EVLA is scheduled on a semester basis, with each semester lasting six months. Proposal deadlines will be 5pm (1700) Eastern Time on February 1 and August 1 (if the deadline falls on a holiday or weekend, it is extended to the next working day). The February 1 proposal deadline nominally covers time to be scheduled during the following August through January, and the August 1 deadline is for time to be scheduled from February through July. Proposals for any configuration in the current '''D'''→'''A''' configuration cycle (September 2011 through January 2013) may be submitted at any proposal deadline, although a proposal for a configuration that has already passed may not be held over for consideration in the next configuration cycle, since the capabilities to be offered in the future are likely to be considerably different from those described in this document. All proposals will be reviewed by a Science Review Panel (SRP) in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The SRP's comments and rating are strongly advisory to the NRAO Time Allocation Committee (TAC), and the comments of both groups are passed on to the proposers soon after each meeting of the TAC (twice yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/observing/ for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2011 and 2012. == Director's Discretionary Time == The NRAO has established two categories of proposals for Director's Discretionary Time (DDT). DDT is limited to a maximum of 5% of the total observing time on the EVLA. All DDT proposals should be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Target of Opportunity.''' Target of Opportunity (ToO) proposals are for unexpected or unpredicted phenomena such as supernovae in nearby galaxies or extreme X-ray or radio flares. ToO Proposals are evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. ToO Proposals are evaluated on the basis of scientific merit by the Chair of the relevant Science Review Panel and Observatory staff with the necessary scientific expertise. The technical feasibility of the proposed observations will be assessed by Observatory staff. The proprietary period for data obtained by ToO Proposals will be assessed on a case-by-case basis but will be no more than six months. 2. '''Exploratory Time.''' Exploratory Proposals are normally for requests of small amounts of time, typically a few hours or less, in response to a recent discovery, possibly to facilitate future submission of a larger proposal. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current EVLA configuration rather than waiting 16 months. The possibility that a proposer forgets about or misses a proposal deadline, or just discovered that he/she was granted time for a particular source on some other telescope, will not constitute sufficient justification for granting observing time by this process. Thus, Exploratory Proposals must include a clear description of why the proposal could not have been submitted for normal review at a previous NRAO proposal deadline, and why it should not wait for the next proposal deadline. Proposals for exploratory time will be evaluated on the basis of scientific merit by the relevant Science Review Panel. Observatory staff will assess their technical feasibility. Notification of the disposition of an Exploratory Proposal normally will be within three weeks of receipt of the proposal; some of these proposals may be put in a queue such that they may or may not be observed. The proprietary period for data obtained by Exploratory Proposals normally will be six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of two formats: :– As a CASA Measurement Set. :– In UVFITS format, which can be read by either AIPS or CASA. The raw SDM format will only be available by special request. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. See http://casa.nrao.edu for more information on the latest release. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aips.nrao.edu for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/nsf06316/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Student Observing Support Program == In addition to travel support for individual data reduction visits NRAO maintains a program to support research by students, both graduate and undergraduate, at U.S. universities and colleges. Regular and Large proposals submitted for the EVLA, VLBA, and GBT, and any combination of these telescopes, are eligible. New applications to the program may be submitted along with new observing proposals at any proposal deadline. Details of this program can be found at https://science.nrao.edu/opportunities/student-programs/studentprograms == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aips.nrao.edu/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aips.nrao.edu/goaips.html. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa_cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS; data reduction and imaging algorithms |- | Miriam Krauss || 7230 || 300 || EVLA CASA subsystem scientist; rapid response science |- | Chris Langley || 7145 || 328 || EVLA Project Manager |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning; EVLA user support |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || WIDAR subsystem scientist; EVLA scientific software |- | Debra Shepherd || 7315 || 330 || EVLA Commissioning |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: queries should generally be directed to the NRAO Helpdesk, at http://science.nrao.edu/observing/helpdesk.shtml. However, you may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the editors of the present document (gvanmoor at nrao dot edu, cchandle at nrao dot edu) with questions on the material, or suggestions that would enhance the clarity of this guide. f73fdfc39967054b04e9fd8951470bdb385be4e8 1439 1438 2013-07-17T19:55:42Z Gvanmoor 7 wikitext text/x-wiki '''REPLACED ON 23 MAY 2012. FOR AN OVERVIEW OF CURRENT AND PAST VERSIONS SEE https://science.nrao.edu/facilities/vla/oss/oss''' '''WE RECOMMEND YOU UPDATE YOUR BOOKMARKS ACCORDINGLY''' ---- '''The EVLA Observational Status Summary''' ''Version date: November 29, 2011'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget at the end of 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || align='center'| 10 μJy || align='center'| 1 μJy || align='center'| 10 |- | Maximum BW in each polarization || align='center'| 0.1 GHz || align='center'| 8 GHz || align='center'| 80 |- | Number of frequency channels at max. BW || align='center'| 16 || align='center'| 16,384 || align='center'| 1024 |- | Maximum number of freq. channels || align='center'| 512 || align='center'| 4,194,304 || align='center'| 8192 |- | Coarsest frequency resolution || align='center'| 50 MHz || align='center'| 2 MHz || align='center'| 25 |- | Finest frequency resolution || align='center'| 381 Hz || align='center'| 0.12 Hz || align='center'| 3180 |- | Number of full-polarization sub-correlators || align='center'| 2 || align='center'| 64 || align='center'| 32 |- | Log (Frequency Coverage over 1–50 GHz) || align='center'| 22% || align='center'| 100% || align='center'| 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date !Status or Actual Date |- | Installation of EVLA correlator subset for early science || align='center'| 2010 Q1 || align='center' | 2012 Q1 |- | Shared Risk Observing begins || align='center'| 2010 Q1 || align='center' | 2012 Q2 |- | Last antenna retrofitted || align='center'| 2010 Q2 || align='center' | 2010 Q2 |- | Full EVLA correlator installation || align='center'| 2011 Q2 || align='center' | 2011 Q2 |- | Last receiver installed || align='center'| 2012 Q4 || align='center' | on schedule |- |} == VLA to EVLA Transition == The year 2010 was extremely exciting for the EVLA. The correlator that was the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from '''A'''→'''B'''→'''C'''→'''D'''→'''A''' to '''D'''→'''C'''→'''B'''→'''A'''→'''D''', in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. The last VLA antenna was retrofitted to EVLA specifications in May 2010. During 2011 the WIDAR correlator was put into full observing mode with commissioning and Resident Shared Risk Observing starting in early 2011. By the end of 2011, up to 2 GHz of bandwidth was provided to the the general public along with 2 tunable bands, each with 8 spectral windows (64 to 256 channels) at S, Ku and X bands. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it has not yet been commissioned. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted '''D''', '''C''', '''B''', and '''A''' respectively. In addition, there are 3 “hybrid” configurations labelled '''DnC''', '''CnB''', and '''BnA''', in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle was modified during 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule for 2011, 2012 and 2013 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Approximate EVLA Configuration Schedule for 2011-2012''' ! Year ! Jan-Apr ! May-Aug ! Sep-Dec |- | align='center'| 2011 || align='center'| '''B''' || align='center'| '''A''' || align='center'| '''D''' |- | align='center'| 2012 || align='center'| '''C''' || align='center'| '''B''' || align='center'| '''A''' |- | align='center'| 2013 || align='center'| '''D''' || align='center'| '''C''' || align='center'| '''B''' |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 2012 (a full '''D'''→'''A''' configuration cycle) are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science is provided by two programs for outside users and one for EVLA commissioning staff. All early science programs are peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period thus involves an element of risk associated with the large stepwise increases in throughput bandwidth that are offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program provides early science capabilities to the general user community. These capabilities initially provided a maximum 256 MHz bandwidth that increased to 2 GHz for the '''D''' configuration in mid-2011 and will increase further to 8 GHz at the end of 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All EVLA antennas are outfitted with either EVLA or “interim” L, EVLA C, VLA X, EVLA K, EVLA Ka, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of July 2012, 25 of the EVLA antennas will be outfitted with S-band receivers, 23 EVLA antennas will have Ku-band receivers, and 22 will have new EVLA-style X-band receivers. Figure 1 shows the expected installation rate of final EVLA receiver systems for the rest of the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers, expected to be commissioned by the end of 2012. Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:RcvrAvailDec12.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of final EVLA receivers during the calendar year 2012 until the end of the EVLA Construction Project on December 31, 2012. Interim receivers with reduced frequency coverage or polarization purity are available at some bands in addition to those plotted (see Table 4).]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that are available in September 2011 and January 2012, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, Jan 2012 ! colspan="2" | Receiver availability, Jul 2012 |- ! ! GHz ! Final EVLA ! EVLA+VLA/interim ! Final EVLA ! EVLA+VLA/interim |- | 400 cm (4-band) || align='center'| 0.062–0.078 || align='center'| - || align='center'| - || align='center'|- || align='center'|- |- | 20 cm (L) || align='center'| 1.0–2.0 || align='center'| 19 || align='center'| 27 || align='center'|24 || align='center'|27 |- | 13 cm (S) || align='center'| 2.0–4.0 || align='center'| 21 || align='center'| 21 || align='center'|25 || align='center'|25 |- | 6 cm (C) || align='center'| 4.0–8.0 || align='center'| 27 || align='center'| 27 || align='center'|27 || align='center'|27 |- | 3 cm (X) || align='center'| 8.0–12.0 || align='center'| 16 || align='center'| 27 || align='center'|22 || align='center'|27 |- | 2 cm (Ku) || align='center'| 12.0–18.0 || align='center'| 18 || align='center'| 18 || align='center'|23 || align='center'|23 |- | 1.3 cm (K) || align='center'| 18.0–26.5 || align='center'| 27 || align='center'| 27 || align='center'|27 || align='center'|27 |- | 1 cm (Ka) || align='center'| 26.5–40.0 || align='center'| 27 || align='center'| 27 || align='center'|27 || align='center'|27 |- | 0.7 cm (Q) || align='center'| 40.0–50.0 || align='center'| 27 || align='center'| 27 || align='center'|27 || align='center'|27 |- |} :Note: The "EVLA+VLA/interim" columns give the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers (see Figure 1). == Open Shared Risk Observing (OSRO) == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program has extended this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program obtain good quality data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing (RSRO) == The WIDAR correlator and the EVLA provides a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program is now expected to run through the end of 2012, with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. Full operations of the EVLA will begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program for the upcoming '''D'''→'''A''' configuration cycle, Sep 2011 through Dec 2012. Users interested in participating in the RSRO program should refer to the web page at https://science.nrao.edu/facilities/evla/early-science/rsro for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24 || 80 || 260 || 850 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka) || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800 || 2200 || 20000 || 20000 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka) || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations ('''DnC''', '''CnB''', '''BnA''') should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., '''DnB''', or '''CnA''') is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. For the '''C''' configuration an antenna from the middle of the north arm is moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Note that although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration. ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are insufficient EVLA 8–12 GHz receivers available yet to determine system performance across the band. A project with the goal of doubling the longest baseline available in the '''A''' configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the ''SEFD'' at some fiducial EVLA frequencies. [[File:SEFD.png|none|frame|Figure 2: ''SEFD'' for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the Ku, K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: ''SEFD''s and '''D'''-Configuration Confusion Limits''' !Frequency !''SEFD'' !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || align='center'| 420 || align='center'| 89 |- | 3.0 GHz (S) || align='center'| 370 || align='center'| 14 |- | 6.0 GHz (C) || align='center'| 310 || align='center'| 2.3 |- | 10.0 GHz (X) || align='center'| 250 || align='center'| negligible |- | 15 GHz (Ku) || align='center'| 350 || align='center'| negligible |- | 22 GHz (K) || align='center'| 560 || align='center'| negligible |- | 33 GHz (Ka) || align='center'| 730 || align='center'| negligible |- | 45 GHz (Q) || align='center'| 1400 || align='center'| negligible |} :Note: ''SEFD''s at Ku, K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in '''D''' configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temperature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted above assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. '''At source elevations greater than 80 degrees (zenith angle < 10 degrees), source tracking becomes difficult; it is recommended to avoid such source elevations during the observation preparation setup'''. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in Right Ascension (RA). The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where ''T''_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and ''S'' in mJy per beam, the constant ''F'' depends only upon array configuration and has the approximate value ''F'' = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for ''S''. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of ''F'' calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where ''D'' is the distance to the galaxy in Mpc, and ''S∆V'' is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different “baseband” frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” Currently, a maximum of 1.024 GHz can be correlated for each IF pair (see [[#Correlator Configurations|Correlator Configurations]]), for a total maximum bandwidth of approximately 2 GHz. The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. At X-band a number of the antennas will continue to have the old narrow-band VLA X-band receivers until their retrofit is complete at the end of the EVLA construction project. As of December 2010 there is not a sufficient number of EVLA-style X-band receivers in the array to evaluate either the system performance or the radio frequency interference environment throughout the EVLA X-band tuning range of 8-12 GHz. A total bandwidth of 800 MHz equivalent to that of the VLA receivers (8.0-8.8 GHz) should be assumed for the purposes of sensitivity calculations at X-band for the present. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || align='center'| 1.0–2.0 || align='center'| 1.25¹ || align='center'| 1.75¹ |- | 13 cm (S) || align='center'| 2.0–4.0 || align='center'| 2.5 || align='center'| 3.5 |- | 6 cm (C) || align='center'| 4.0–8.0 || align='center'| 5.0 || align='center'| 6.0 |- | 3 cm (X) || align='center'| 8.0–12.0 || align='center'| 8.5² || align='center'| 9.5² |- | 2 cm (Ku) || align='center'| 12.0–18.0 || align='center'| 13.5 || align='center'| 14.5 |- | 1.3 cm (K) || align='center'| 18.0–26.5 || align='center'| 20.7 || align='center'| 21.7 |- | 1 cm (Ka) || align='center'| 26.5–40.0 || align='center'| 31.5 || align='center'| 32.5 |- | 0.7 cm (Q) || align='center'| 40.0–50.0 || align='center'| 40.5 || align='center'| 41.5 |} :Notes: :: 1. This default frequency set-up for L-band comprises two 512 MHz basebands (each with 8 subbands of 64 MHz) to cover the entire 1-2 GHz of the L-band receiver. :: 2. Many of the antennas continue to have the old narrow-band VLA receivers (see Figure 1), for which a total bandwidth of 800 MHz should be assumed (8.0-8.8 GHz). The RFI environment of the default tunings has not yet been evaluated. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 in [[#Documentation|Documentation]] for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! (∆ν/ν₀)*(θ₀/θ_HPBW) ! Peak ! Width |- | 0.0 || align='center'| 1.00 || align='center'| 1.00 |- | 0.50 || align='center'| 0.95 || align='center'| 1.05 |- | 0.75 || align='center'| 0.90 || align='center'| 1.11 |- | 1.0 || align='center'| 0.80 || align='center'| 1.25 |- | 2.0 || align='center'| 0.50 || align='center'| 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 (in [[#Documentation|Documentation]]) considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || align='center'| 2.1 || align='center'| 4.8 || align='center'| 6.7 |- | '''B''' || align='center'| 6.8 || align='center'| 15.0 || align='center'| 21.0 |- | '''C''' || align='center'| 21.0 || align='center'| 48.0 || align='center'| 67.0 |- | '''D''' || align='center'| 68.0 || align='center'| 150.0 || align='center'| 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 (http://www.aoc.nrao.edu/evla/geninfo/memoseries/evlamemo67.pdf) for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second, for '''A''' configuration. For '''B''', '''C''', and '''D''' configurations the minimum integration time is 3 seconds, unless a shorter integration time is explicitly requested and justified. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 14.4 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 51.8 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 1.2 TB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool (https://science.nrao.edu/facilities/evla/data-archive/evla) will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on disk drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the '''D''' configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum at L-band. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots and tables of known RFI are available online, at http://science.nrao.edu/evla/observing/RFI/; plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the '''D''' configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in '''A''' configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in '''A''' configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List (http://www.vla.nrao.edu/astro/calib/manual/) should be used. The positions of these sources are taken from lists published by the United States Naval Observatory (USNO). == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1 in [[#Documentation|Documentation]]. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth smearing and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its '''D''' configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP (Wilkinson Microwave Anisotropy Probe) flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby calibrator with an "S" code in the calibrator database, and a more distant calibrator with a "P" code, the nearby calibrator is usually the better choice (see http://www.vla.nrao.edu/astro/calib/manual/key.html for a description of calibrator codes). :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173 (http://www.vla.nrao.edu/memos/sci/). These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. A detailed description of the API may be found at http://www.vla.nrao.edu/astro/guides/api/. Plots of current/historical data can be found at: https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi Characteristic seasonal averages are represented below: {| border="1" align="center" |+ '''Table: Seasonal API/wind values at the EVLA''' !Month !API (night) [deg] !API (median) [deg] !API (day) [deg] !Wind (night) [m/s] !Wind (median) [m/s] !Wind (day) [m/s] |- | [[Media:APIwind_January.png| January]] || 2.3 || 2.8 || 3.6 || 1.6 || 1.9 || 2.3 |- | [[Media:APIwind_February.png| February]] || 2.9 || 3.4 || 4.5 || 4.0 || 4.3 || 4.5 |- | [[Media:APIwind_March.png| March]] || 2.8 || 3.7 || 5.5 || 3.4 || 3.9 || 4.7 |- | [[Media:APIwind_April.png| April]] || 3.3 || 4.5 || 6.2 || 5.3 || 5.5 || 5.8 |- | [[Media:APIwind_May.png | May]] || 2.9 || 4.6 || 6.7 || 2.6 || 3.2 || 3.7 |- | [[Media:APIwind_June.png| June]] || 3.8 || 5.5 || 7.4 || 2.5 || 3.9 || 6.3 |- | [[Media:APIwind_July.png| July]] || 6.2 || 8.3 || 10.5 || 2.9 || 2.9 || 3.0 |- | [[Media:APIwind_August.png| August]] || 5.4 || 7.1 || 11.3 || 1.7 || 2.3 || 3.0 |- | [[Media:APIwind_September.png| September]] || 5.2 || 6.6 || 8.8 || 2.3 || 3.0 || 3.6 |- | [[Media:APIwind_October.png| October]] || 4.2 || 5.3 || 7.4 || 2.3 || 2.9 || 3.7 |- | [[Media:APIwind_November.png| November]] || 2.6 || 3.0 || 4.0 || 1.2 || 2.5 || 1.6 |- | [[Media:APIwind_December.png| December]] || 2.8 || 3.2 || 4.1 || 1.2 || 1.6 || 2.7 |- |} Click on the Month links above to see plots of phase and wind speed versus time. :Note: day indicates sunrise to sunset values; night indicates sunset to sunrise values. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.aoc.nrao.edu/~smyers/evlapolcal/polcal_master.html High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/<math>{\rm{m}^2}</math>, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. For more details see: https://science.nrao.edu/facilities/evla/early-science/polarimetry == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. For the Open Shared Risk Observing (OSRO) program available to the community during the period Sep 2011 through Jan 2013 are offering two independently tunable basebands, where each baseband has up to eight sub-bands. Possible sub-band widths are 128 MHz, 64 MHz, 32 MHz, all the way down in factors of 2 to 0.03125 MHz. All sub-bands must have the same bandwidth and channelization in both basebands, and be contiguous in frequency within each baseband. We will offer three different OSRO modes: full polarization, dual polarization, and single polarization, with 64, 128, and 256 channels per sub-band, respectively. There is always the possibility during offline processing to smooth in frequency to reduce dataset sizes or to improve spectral response. Starting with the '''D'''-configuration in September 2011, we have been providing options for configuring WIDAR for OSRO in the following three ways: :1. “OSRO Full Polarization”: Four polarization products. This configuration offers 4 polarization products for each sub-band, each of which has 128 MHz bandwidth with 64 channels. It is possible to decrease the subband bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in the following table: {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO Full Polarization)''' ! Sub-band BW (MHz) ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 64 || 1000 || 300 || 19,200 |- | 32 || 64 || 500 || 150 || 9,600 |- | 16 || 64 || 250 || 75 || 4,800 |- | 8 || 64 || 125 || 37.5 || 2,400 |- | 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 64 || 31.25 || 9.4 || 600 |- | 1 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO Dual Polarization”: Two polarization products. This configuration offers 2 polarization products for each sub-band, each of which has 128 MHz bandwidth with 128 channels. It is possible to decrease the sub-band bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in the following table. {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for dual polarization (OSRO Dual Polarization)''' ! Sub-band BW (MHz) ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 128 || 1000 || 300/ν (GHz) || 38,400/ν (GHz) |- | 64 || 128 || 500 || 150 || 19,200 |- | 32 || 128 || 250 || 75 || 9,600 |- | 16 || 128 || 125 || 37.5 || 4,800 |- | 8 || 128 || 62.5 || 19 || 2,400 |- | 4 || 128 || 31.25 || 9.4 || 1,200 |- | 2 || 128 || 15.625 || 4.7 || 600 |- | 1 || 128 || 7.813 || 2.3 || 300 |- | 0.5 || 128 || 3.906 || 1.2 || 150 |- | 0.25 || 128 || 1.953 || 0.59 || 75 |- | 0.125 || 128 || 0.977 || 0.29 || 37.5 |- | 0.0625 || 128 || 0.488 || 0.15 || 18.75 |- | 0.03125 || 128 || 0.244 || 0.073 || 9.375 |} :3: "OSRO Single Polarization": One polarization product (new for OSRO observing). It offers 1 polarization product for each sub-band, each of which has 128 MHz bandwidth with 256 channels. It is possible to decrease the sub-band bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in the following table. {| border="1" align="center" |+ '''Table 14: Correlator capabilities per sub-band for single polarization (OSRO Single Polarization)''' ! Sub-band BW (MHz) ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 256 || 250 || 75 || 19,200 |- | 32 || 256 || 125 || 37.5 || 9,600 |- | 16 || 256 || 62.5 || 19 || 4,800 |- | 8 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities are being provided with integration times no shorter than 1 second in '''A''' configuration (3 seconds in '''B'''/'''C'''/'''D''' configurations), and Doppler setting will be available with these correlator configurations. If it is likely that the data will need to be resampled spectrally in order to Doppler track to a line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna (e.g., Y1) modes, have not yet been commissioned and are not yet available to the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing "Faint Images of the Radio Sky at Twenty-centimeters survey(FIRST, http://www.cv.nrao.edu/first/) or the Co-Ordinated Radio 'N' Infrared Survey for High-mass star formation (CORNISH, http://www.ast.leeds.ac.uk/Cornish/public/index.php) ('''B''' configuration), or the NRA VLA Sky Survey (NVSS, http://www.cv.nrao.edu/nvss/) ('''D''' configuration, all-sky) surveys. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 ([[#Documentation|Documentation]]) for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to obtain EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may find that such dissertations comprise pieces of several short proposals, which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being impaired by an adverse review of one proposal when the full scope of the project is not seen. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. Starting in 2011 time on the EVLA is scheduled on a semester basis, with each semester lasting six months. Proposal deadlines will be 5pm (1700) Eastern Time on February 1 and August 1 (if the deadline falls on a holiday or weekend, it is extended to the next working day). The February 1 proposal deadline nominally covers time to be scheduled during the following August through January, and the August 1 deadline is for time to be scheduled from February through July. Proposals for any configuration in the current '''D'''→'''A''' configuration cycle (September 2011 through January 2013) may be submitted at any proposal deadline, although a proposal for a configuration that has already passed may not be held over for consideration in the next configuration cycle, since the capabilities to be offered in the future are likely to be considerably different from those described in this document. All proposals will be reviewed by a Science Review Panel (SRP) in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The SRP's comments and rating are strongly advisory to the NRAO Time Allocation Committee (TAC), and the comments of both groups are passed on to the proposers soon after each meeting of the TAC (twice yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/observing/ for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2011 and 2012. == Director's Discretionary Time == The NRAO has established two categories of proposals for Director's Discretionary Time (DDT). DDT is limited to a maximum of 5% of the total observing time on the EVLA. All DDT proposals should be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Target of Opportunity.''' Target of Opportunity (ToO) proposals are for unexpected or unpredicted phenomena such as supernovae in nearby galaxies or extreme X-ray or radio flares. ToO Proposals are evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. ToO Proposals are evaluated on the basis of scientific merit by the Chair of the relevant Science Review Panel and Observatory staff with the necessary scientific expertise. The technical feasibility of the proposed observations will be assessed by Observatory staff. The proprietary period for data obtained by ToO Proposals will be assessed on a case-by-case basis but will be no more than six months. 2. '''Exploratory Time.''' Exploratory Proposals are normally for requests of small amounts of time, typically a few hours or less, in response to a recent discovery, possibly to facilitate future submission of a larger proposal. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current EVLA configuration rather than waiting 16 months. The possibility that a proposer forgets about or misses a proposal deadline, or just discovered that he/she was granted time for a particular source on some other telescope, will not constitute sufficient justification for granting observing time by this process. Thus, Exploratory Proposals must include a clear description of why the proposal could not have been submitted for normal review at a previous NRAO proposal deadline, and why it should not wait for the next proposal deadline. Proposals for exploratory time will be evaluated on the basis of scientific merit by the relevant Science Review Panel. Observatory staff will assess their technical feasibility. Notification of the disposition of an Exploratory Proposal normally will be within three weeks of receipt of the proposal; some of these proposals may be put in a queue such that they may or may not be observed. The proprietary period for data obtained by Exploratory Proposals normally will be six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of two formats: :– As a CASA Measurement Set. :– In UVFITS format, which can be read by either AIPS or CASA. The raw SDM format will only be available by special request. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. See http://casa.nrao.edu for more information on the latest release. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aips.nrao.edu for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/nsf06316/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Student Observing Support Program == In addition to travel support for individual data reduction visits NRAO maintains a program to support research by students, both graduate and undergraduate, at U.S. universities and colleges. Regular and Large proposals submitted for the EVLA, VLBA, and GBT, and any combination of these telescopes, are eligible. New applications to the program may be submitted along with new observing proposals at any proposal deadline. Details of this program can be found at https://science.nrao.edu/opportunities/student-programs/studentprograms == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aips.nrao.edu/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aips.nrao.edu/goaips.html. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa_cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS; data reduction and imaging algorithms |- | Miriam Krauss || 7230 || 300 || EVLA CASA subsystem scientist; rapid response science |- | Chris Langley || 7145 || 328 || EVLA Project Manager |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning; EVLA user support |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || WIDAR subsystem scientist; EVLA scientific software |- | Debra Shepherd || 7315 || 330 || EVLA Commissioning |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: queries should generally be directed to the NRAO Helpdesk, at http://science.nrao.edu/observing/helpdesk.shtml. However, you may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the editors of the present document (gvanmoor at nrao dot edu, cchandle at nrao dot edu) with questions on the material, or suggestions that would enhance the clarity of this guide. d73f9635d0e7f28878481af27fb8f8107f9cdf3a 0 145 1441 2013-07-17T19:57:03Z Gvanmoor 7 wikitext text/x-wiki '''WARNING: OBSOLETE CONTENT. FOR AN OVERVIEW OF CURRENT AND PAST VERSIONS SEE https://science.nrao.edu/facilities/vla/oss/oss''' '''WE RECOMMEND YOU UPDATE YOUR BOOKMARKS ACCORDINGLY''' '''The EVLA Observational Status Summary''' ''Version date: May 12, 2010'' = Introduction = == Purpose of Document == This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in [[#Key Personnel|Key Personnel]], or refer to the manuals and documentation listed in [[#Documentation|Documentation]]. Most of the information contained here, and much more, is available through the EVLA science web pages, at http://science.nrao.edu/evla/. A companion document for the VLBA is available at http://science.nrao.edu/vlba/. The EVLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it. Although the NRAO strives to make using the EVLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to [[#Help for Visitors to the EVLA and DSOC|Help for Visitors to the EVLA and DSOC]] for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it. == What is the Expanded Very Large Array? == The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is approximately $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $2 million from Mexico. The EVLA project will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2. {| border="1" align="center" |+ '''Table 1: Overall EVLA Performance Goals''' !Parameter !VLA !EVLA !Factor |- | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10 |- | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80 |- | Number of frequency channels at max. BW || 16 || 16,384 || 1024 |- | Maximum number of freq. channels || 512 || 4,194,304 || 8192 |- | Coarsest frequency resolution || 50 MHz || 2 MHz || 25 |- | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180 |- | Number of full-polarization sub-correlators || 2 || 64 || 32 |- | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5 |- |} :Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter. {| border="1" align="center" |+ '''Table 2: EVLA Major Milestones''' !Milestone !Target Date |- | Installation of EVLA correlator subset for early science || 2010 Q1 |- | Shared Risk Observing begins || 2010 Q1 |- | Last antenna retrofitted || 2010 Q2 |- | Full EVLA correlator installation || 2010 Q3 |- | Last receiver installed || 2012 Q4 |- |} == VLA to EVLA Transition == The year 2010 is extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA “WIDAR” correlator. The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. = An Overview of the EVLA = The EVLA is a 27-element interferometric array, arranged along the arms of an upside-down “Y”, which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico. The basic data produced by the EVLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array’s elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing (commonly called phased array) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially. The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A respectively. In addition, there are 3 “hybrid” configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15◦ or north of δ = +75◦, for which the foreshortening of the longer North arm results in a more circular point spread function. Traditionally, the VLA completed one cycle through all four configurations in an approximately 16 month period. However, the length of the cycle may in 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to [[#Obtaining Observing Time on the EVLA|Obtaining Observing Time on the EVLA]] for information on how to submit an observing proposal. {| border="1" align='center' |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011''' ! Year ! Mar-Sep ! Oct-Jan |- | 2010 || '''D''' || '''C''' |- ! Year ! Feb-Apr ! May-Aug ! Sep-Dec |- | 2011 || '''B''' || '''A''' || '''D''' |- |} Observing projects on the EVLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources. However, since the EVLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See [[#Snapshots|Snapshots]] for details. All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 1–2 GHz, 2–4 GHz, 4–8 GHz, 8–12 GHz, 12–18 GHz, 18–26.5 GHz, 26.5–40 GHz, and 40–50 GHz. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See [[#Expected Capabilities|Expected Capabilities]] for more details about the availability of new bands. The VLA’s original P-band (300–340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will eventually also replace the existing 74-MHz (4-band) receivers. In the interim, a set of six 74-MHz dipoles were mounted with the wideband EVLA electronics in the DnC-configuration, providing information on the low frequency observing capability (response over 16 MHz BW in the 62-78 MHz range, with a significant roll-off in the low frequency end due to the dipole response); the October 1 2010 call for proposals will accept observations in this band for the B/BnA/A configurations; the sensitivity should be assumed to be that of the old VLA system (RMS (10 min)=160 mJy), though the improved correlator, bandwidth and interference environment may significantly improve this value (See: http://evlaguides.nrao.edu/index.php?title=File:4-bandResponse.png). The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in [[#Correlator Configurations|Correlator Configurations]]. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult [[#Correlator Configurations|Correlator Configurations]]. = EVLA Early Science = EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012. The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program. == Expected Capabilities == All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of the beginning of May 2010, 21 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 8 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization. [[File:WideBandRcvrFrcstMay10.png|none|frame|Figure 1: EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.]] Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in May 2011, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff. {| border="1" align="center" |+ '''Table 4: Tuning Ranges of EVLA Bands''' ! Band ! Range ! colspan="2" | Receiver availability, May 2011 |- ! ! GHz ! Final EVLA systems ! Total EVLA+VLA/interim |- | 400 cm (4-band) || 0.062-0.078 || || align='center'| 27 |- | 20 cm (L) || 1.0-2.0 || align='center'| 16 || align='center'| 27 |- | 13 cm (S) || 2.0-4.0 || align='center'| 16 || align='center'| 16 |- | 6 cm (C) || 4.0-8.0 || align='center'| 27 || align='center'| 27 |- | 3 cm (X) || 8.0-12.0 || align='center'| 7 || align='center'| 27 |- | 2 cm (Ku) || 12.0-18.0 || align='center'| 11 || align='center'| 11 |- | 1.3 cm (K) || 18.0-26.5 || align='center'| 27 || align='center'| 27 |- | 1 cm (Ka) || 26.5-40.0 || align='center'| 27 || align='center'| 27 |- | 0.7 cm (Q) || 40.0-50.0 || align='center'| 27 || align='center'| 27 |- |} :Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers (see Figure 1). == Open Shared Risk Observing == NRAO has been offering “shared risk” access to the VLA for all users since the EVLA construction project began. The OSRO program extends this into the EVLA era by providing early access to a number of EVLA correlator capabilities and observing modes that represent a significant improvement over the capabilities of the VLA correlator. They are described in detail in [[#Performance of the EVLA|Performance of the EVLA]]. NRAO will make every effort to ensure projects awarded time under the OSRO program do obtain data. The highest risk will be for time-critical observations such as observations of triggered transients or observations coordinated with other observatories. == Resident Shared Risk Observing == The WIDAR correlator and the EVLA will provide a vastly more powerful instrument than the VLA. The RSRO program offers participants early access to the growing capabilities of the EVLA as it is being commissioned, in exchange for a period of residence in Socorro to assist with the commissioning. It is intended to accelerate the development of the EVLA’s full scientific capabilities by gaining enhanced resources and expertise through community participation. It will at the same time help quickly optimize the scientific productivity of the EVLA. The RSRO program will run for approximately two years (from March 2010 through the end of 2011), with up to 25% of the EVLA time available for astronomical observations allocated to the RSRO program, depending on demand and quality of science proposed. At the end of this period all access to the EVLA will be through the OSRO program, until full operations begin in 2013. This document describes only those capabilities available to the general user community through the OSRO program. Users interested in participating in the RSRO program should refer to the web page at http://science.nrao.edu/evla/earlyscience/rsro.shtml for a general description of the expected capabilities available through the RSRO program, and should check with EVLA staff for further details. = Performance of the EVLA = This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in [[#Documentation|Documentation]]]. == Resolution == The EVLA’s resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is “blind” to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information. Table 5 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power), and the scale at which severe attenuation of large scale structure occurs. {| border="1" align="center" |+ '''Table 5: Configuration Properties''' !Configuration !A !B !C !D |- |B_max (km¹) || 36.4 || 11.1 || 3.4 || 1.03 |- |B_min (km¹) || 0.68 || 0.21 || 0.035⁵ || 0.035 |- | ||colspan="4"| Synthesized Beamwidth θ_HPBW(arcsec)^1,2,3 |- |74 MHz (4 band) || 24 || 80 || 260 || 850 |- |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46 |- |3.0 GHz (S)⁶ || 0.65 || 2.1 || 7.0 || 23 |- |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12 |- |8.5 GHz (X)⁷ || 0.23 || 0.73 || 2.5 || 8.1 |- |15 GHz (Ku)⁶ || 0.13 || 0.42 || 1.4 || 4.6 |- |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1 |- |33 GHz (Ka)⁶ || 0.059 || 0.19 || 0.63 || 2.1 |- |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5 |- | || colspan="4" | Largest Angular Scale θ_LAS(arcsec)^1,4 |- |74 MHz (4 band) || 800 || 2200 || 20000 || 20000 |- |1.5 GHz (L) || 36 || 120 || 970 || 970 |- |3.0 GHz (S)⁶ || 18 || 58 || 490 || 490 |- |6.0 GHz (C) || 8.9 || 29 || 240 || 240 |- |8.5 GHz (X)⁷ || 6.3 || 20 || 170 || 170 |- |15 GHz (Ku)⁶ || 3.6 || 12 || 97 || 97 |- |22 GHz (K) || 2.4 || 7.9 || 66 || 66 |- |33 GHz (Ka)⁶ || 1.6 || 5.3 || 44 || 44 |- |45 GHz (Q) || 1.2 || 3.9 || 32 || 32 |- |} :These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith. :Footnotes: ::1. B_max is the maximum antenna separation, B_min is the minimum antenna separation, θ_HPBW is the synthesized beam width (FWHM), and θ_LAS is the largest scale structure "visible" to the array. ::2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (BnA, CnB, DnC) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., CnA, or DnB) is provided. ::3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θ_HPBW. The values for snapshots are about 1.3 times the listed values. ::4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two. ::5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration). ::6. The S and Ku bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations. ::7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010. A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed. == Sensitivity == The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by: <math>\Delta I_m = \frac{SEFD}{\eta_{\rm c}\sqrt{n_{\rm pol}N(N-1)t_{\rm int}\Delta\nu}}</math> ::::: ...equation (1) where: :– ''SEFD'' is the “system equivalent flux density” (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the ''SEFD'' indicate more sensitive performance. For the EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T_sys/η_A, where T_sys is the total system temperature (receiver plus antenna plus sky), and η_A is the antenna aperture efficiency in the given band. :– η_c is the correlator efficiency (at least 0.92 for the EVLA). :– n_pol is the number of polarization products included in the image; n_pol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’. :– N is the number of antennas. :– t_int is the total on-source integration time in seconds. :– ∆ν is the bandwidth in Hz. Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. [[File:SEFD.png|none|frame|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]] {| border="1" align="center" |+ '''Table 6: SEFDs and D-Configuration Confusion Limits''' !Frequency !SEFD !RMS confusion level |- ! ! (Jy) ! in '''D''' config (μJy/beam) |- | 1.5 GHz (L) || 420 || 89 |- | 3.0 GHz (S) || 370 || 14 |- | 6.0 GHz (C) || 310 || 2.3 |- | 22 GHz (K) || 560 || ... |- | 33 GHz (Ka) || 730 || ... |- | 45 GHz (Q) || 1400 || ... |} :Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above. No current estimates are available for 4-band (74 MHz) observations but will be established during October 2010 DnC configuration testing. Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical: :• For the more commonly-used “robust” weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values. :• Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 6 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 6; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations. :• Weather. The sky and ground temperature contributions to the total system temper- ature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O₂ transitions at Q-band. :In general, the zenith atmospheric opacity to microwave radiation is very low – typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O₂, is considerably less variable. :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz. :Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item. :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target. :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz. :To achieve better pointing, “referenced pointing” is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 – 3 arcseconds if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution. To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations. The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by: <math>T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S</math> ::::: ...equation (2) where T_b is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/B_max), and S in mJy per beam, the constant F depends only upon array configuration and has the approximate value F = 190, 18, 1.7, 0.16 for '''A''', '''B''', '''C''', and '''D''' configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ∆I_m, for S. Note that Equation 2 is a ''beam-averaged'' surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in [[#Documentation|Documentation]]. For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791): <math>M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot</math> ::::: ...equation (3) where D is the distance to the galaxy in M_pc, and S∆V is the HI line area in units of Jy km/s. == EVLA Frequency Bands and Tunability == For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.” The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used. {| border="1" align="center" |+ '''Table 7: Default frequencies for “continuum” applications''' ! Band ! Range ! colspan="2"| Default frequencies for continuum applications (GHz) |- ! ! (GHz) ! IFs A/C ! IFs B/D |- | 20 cm (L) || 1.0-2.0 || 1.328¹ || 1.456¹ |- | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212 |- | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024 |- | 3 cm (X) || 8.0-12.0 || 8.396² || 8.524² |- | 2 cm (Ku) || 12.0-18.0 || ... || ... |- | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524 |- | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624 |- | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344 |} :Notes: :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference-free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1820 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs. :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers. In general, for all frequency bands except Ka, if the total span of the two independent IFs (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For Ka and Q bands (the only two bands where a span greater than 8 GHz is possible), there are special rules: :• At Ka band, the low frequency edge of the AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency. :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency. == Field of View == At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly: === Primary Beam === The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θ_PB = 45/ν_GHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as “mosaicing,” in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 for details. === Chromatic Aberration (Bandwidth Smearing) === The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array. These effects can be parameterized by the product of the fractional bandwidth (∆ν/ν₀) with the source offset in synthesized beamwidths (θ₀/θ_HPBW). Table 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 should be used to determine how much spectral averaging can be tolerated when imaging a particular field. {| border="1" align="center" |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing''' ! ∆ν/ν₀ θ₀/θ_HPBW ! Peak ! Width |- | 0.0 || 1.00 || 1.00 |- | 0.50 || 0.95 || 1.05 |- | 0.75 || 0.90 || 1.11 |- | 1.0 || 0.80 || 1.25 |- | 2.0 || 0.50 || 2.00 |} :Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ₀/θ_HPBW is the source offset from the phase tracking center in units of the synthesized beam. === Time-Averaging Loss === The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90^◦. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with ∆ν/ν₀ replaced by ω_e∆t_int, where ω_e is the Earth’s angular rotation rate, and ∆t_int is the averaging interval. For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 9, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θ∆t_int)², where θ is the distance from the phase center and ∆t_int is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see Section 4.5) in at least the '''A''' and '''B''' configurations. {| border="1" align="center" |+ '''Table 9: Averaging Time for a Given Amplitude Loss''' ! ! colspan="3"|Amplitude loss |- ! Configuration ! 1.0% ! 5.0% ! 10.0% |- | '''A''' || 2.1 || 4.8 || 6.7 |- | '''B''' || 6.8 || 15.0 || 21.0 |- | '''C''' || 21.0 || 48.0 || 67.0 |- | '''D''' || 68.0 || 150.0 || 210.0 |} :Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines. === Non-Coplanar Baselines === The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D², where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations. Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details. == Time Resolution and Data Rates == The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used. The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), shorter integration times may be preferable. Observers should bear in mind the data rate of the EVLA when planning their observations. For the correlator configurations discussed in [[#Correlator Configurations|Correlator Configurations]], and integration time ∆t, the data rate is: :Data rate = 1.7 MB/sec × N_ant × (N_ant−1)/(27×26)/(∆t/1 sec) ::: = 6.0 GB/hour × N_ant ×(N_ant−1)/(27×26)/(∆t/1 sec) ::::: ...equation (4) which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations. Observers may also request their data on DATs or USB drives. == Radio-Frequency Interference == The bands within the tuning range of the EVLA which are allocated exclusively to radio astronomy are 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands. RFI is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels. Radio frequency interference (RFI) at the EVLA will be an increasing problem to astronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spectrum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures. [[File:Lband_sweep.png|none|frame|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference. Note that the y-axis is in logarithmic units (dB).]] [[File:S-bandRFI.png|none|frame|Figure 4: Spectrum of S-band RFI. This shows the raw spectrum of the lower half of S-Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio. The y-axis is in logarithmic units (dB).]] {| border="1" align="center" |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz''' ! Frequency (MHz) ! Source ! Comments |- | 1025-1150 || Aircraft navigation || Very strong |- | 1200.0 || VLA modem || |- | 1217-1237 || GPS L2 || Very strong |- | 1243-1251 || GLONASS L2 || |- | 1254 || Aeronautical radar || |- | 1263 || Aeronautical radar || |- | 1268 || COMPASS E6 || |- | 1310 || Aeronautical radar || |- | 1317 || Aeronautical radar || |- | 1330 || Aeronautical radar || |- | 1337 || Aeronautical radar || |- | 1376-1386 || GPS L3 || Intermittent |- | 1525-1564 || INMARSAT satellites || |- | 1564-1584 || GPS L1 || Very strong |- | 1598-1609 || GLONASS L1 || |- | 1618-1627 || IRIDIUM satellites || |- | 1642 || 2nd harmonic VLA radios || Sporadic |- | 1683-1687 || GOES weather satellite || |- | 1689-1693 || GOES weather satellite || |- | 1700-1702 || NOAA weather satellite || |- | 1705-1709 || NOAA weather satellite || |- | 1930-1990 || PCS cell phone base stations || |- | 2178-2195 || ??? || |- | 2320-2350 || Satellite radio || |- |} Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu. The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding. The EVLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away. Extracting astonomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish, and subtract RFI signals from interferometer data is ongoing. == Subarrays == The separation of the EVLA into multiple sub-arrays is not currently supported for OSRO. == Positional Accuracy == The accuracy with which an object’s position can be determined is limited by the atmospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]), accuracies of 1 milliarcsecond have been attained with the VLA. If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO. == Limitations on Imaging Performance == Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections. === Image Fidelity === With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits. === Invisible Structures === An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Table 5 gives the largest scale visible to each configuration/band combination. === Poorly Sampled Fourier Plane === Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1. === Sidelobes from Confusing Sources === At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. === Sidelobes from Strong Sources === An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere. == Calibration and the Flux Density Scale == The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/. Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well. Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the EVLA cannot measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 11 shows the flux densities of these sources in January 2010 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band. The flux densities for frequencies below 4 GHz are based on the Baars et al. scale. For frequencies above 4 GHz, the flux densities are based on a model of the emission of Mars, tied to the WMAP flux scale. {| border="1" align="center" |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for January 2010''' ! Source/Frequency (MHz) ! 1465 ! 4885 ! 8435 ! 14965 ! 22460 ! 43340 |- | 3C48 = J0137+3309 || 15.15 || 5.40 || 3.19 || ... || 1.23 || 0.70 |- | 3C138 = J0521+1638 || 8.20 || 4.19 || 2.95 || ... || 1.54 || 1.10 |- | 3C147 = J0542+4951 || 20.68 || 7.65 || 4.56 || ... || 1.81 || 1.08 |- | 3C286 = J1331+3030 || 14.51 || 7.31 || 5.05 || 3.39 || 2.50 || 1.54 |- | 3C295 = J1411+5212 || 21.49 || 6.39 || 3.31 || 1.62 || 0.95 || 0.40 |- | NGC 7027 || 1.495 || 5.36 || 5.80 || ... || 5.40 || 5.35 |- | MARS (Jan 11,2010) || 0.039 || 0.442 || 1.31 || ... || 9.41 || 35.1 |- |} Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for further information. From this work it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable. For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation. == Complex Gain Calibration == === General Guidelines for Gain Calibration === Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below. :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm. :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient. :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale. :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice. :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See [[#Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)|Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)]]. === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) === For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼v_at/2 where v_a is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # [http://www.vla.nrao.edu/memos/sci/169/169.ps 169] and [http://www.vla.nrao.edu/memos/sci/173/173.ps 173]. These memos, and other useful information, can be obtained from Reference 12 in [[#Documentation|Documentation]]. Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms). An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions. Archival API data have been combined with wind speed data to create [[Monthly_Conditions_at_EVLA|monthly plots of observing conditions suitable to each receiver band vs. LST]]. A detailed description of the API hardware and data product may be found in VLA Test Memos [http://www.vla.nrao.edu/memos/test/213/213.pdf 213], [http://www.vla.nrao.edu/memos/test/222/222.pdf 222], and [http://adsabs.harvard.edu/abs/1996PASP..108..441R Radford et al. 1996]. The public [http://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi web interface to the API is available here]. == Polarization == For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization calibrator. The minimum condition that will enable accurate polarization calibration is four observations of a bright source spanning at least 90 degrees in parallactic angle. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used to fix the position angle at wavelengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators can be found at http://www.vla.nrao.edu/astro/calib/polar/. High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA’s new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible. The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing. Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m², so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3–1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements. == Correlator Configurations == The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways: :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it is also possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 4 || 64 || 2000 || 600/ν (GHz) || 38,400/ν (GHz) |- | 64 || 4 || 64 || 1000 || 300 || 19,200 |- | 32 || 4 || 64 || 500 || 150 || 9,600 |- | 16 || 4 || 64 || 250 || 75 || 4,800 |- | 8 || 4 || 64 || 125 || 37.5 || 2,400 |- | 4 || 4 || 64 || 62.5 || 19 || 1,200 |- | 2 || 4 || 64 || 31.25 || 9.4 || 600 |- | 1 || 4 || 64 || 15.625 || 4.7 || 300 |- | 0.5 || 4 || 64 || 7.813 || 2.3 || 150 |- | 0.25 || 4 || 64 || 3.906 || 1.2 || 75 |- | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5 |- | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75 |- | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375 |} :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It is possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see [[#EVLA Frequency Bands and Tunability|EVLA Frequency Bands and Tunability]]). {| border="1" align="center" |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)''' ! Sub-band BW (MHz) ! Number of poln. products ! Number of channels/poln. product ! Channel width (kHz) ! Channel width (km/s at 1 GHz) ! Total velocity coverage per sub-band (km/s at 1 GHz) |- | 128 || 2 || 256 || 500 || 150/ν (GHz) || 38,400/ν (GHz) |- | 64 || 2 || 256 || 250 || 75 || 19,200 |- | 32 || 2 || 256 || 125 || 37.5 || 9,600 |- | 16 || 2 || 256 || 62.5 || 19 || 4,800 |- | 8 || 2 || 256 || 31.25 || 9.4 || 2,400 |- | 4 || 2 || 256 || 15.625 || 4.7 || 1,200 |- | 2 || 2 || 256 || 7.813 || 2.3 || 600 |- | 1 || 2 || 256 || 3.906 || 1.2 || 300 |- | 0.5 || 2 || 256 || 1.953 || 0.59 || 150 |- | 0.25 || 2 || 256 || 0.977 || 0.29 || 75 |- | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5 |- | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75 |- | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375 |} These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations. The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011. If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing. All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator. == VLBI Observations == VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program. == Snapshots == The two-dimensional geometry of the EVLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest. Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm in the '''D''' configuration, for which a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. Use of the AIPS program IMAGR or CASA task clean is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' configuration, all-sky) surveys. These surveys can be accessed from the NRAO website, at http://science.nrao.edu/evla/proposing/largeproposals.shtml. == Shadowing and Cross-Talk == Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference. == Combining Configurations and Mosaicing == Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information. == Pulsar Observing == Observation of pulsars requiring modes other than those described in [[#Correlator Configurations|Correlator Configurations]] will not initially be supported on the EVLA. = Using the EVLA = == Obtaining Observing Time on the EVLA == Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information. It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information. Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations. The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configuration is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted. All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the comments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process. Because of competition, even highly rated proposals are not guaranteed to receive observing time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010. == Rapid Response Science == The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered. 1. '''Known Transient Phenomena.''' These proposals will request time to observe phenomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months. 2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal deadline(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted proposals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months. 3. '''Target of Opportunity.''' These proposals are for true targets of opportunity – unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the proposed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months. == Helpdesk == Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk at http://science.nrao.edu/observing/helpdesk.shtml. == Observation Preparation == To use the EVLA, scheduling blocks must be prepared using the “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml. == Fixed Date and Dynamic Scheduling == Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project. Please see the scheduling officers’ home page for further information, at http://science.nrao.edu/evla/sched/schedsoc.shtml. == The Observations and Remote Observing == Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]] for information on coming to and staying in Socorro. For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format. == Data Access == The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://science.nrao.edu/evla/archive/evla/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see [[#The Observations and Remote Observing|The Observations and Remote Observing]]). Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats: :– SDM format, which must be read into CASA using the task importasdm for further processing. :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm. :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS. == Data Processing == The primary data reduction package for the EVLA is the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information. NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details. It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS. == Travel Support for Visiting the DSOC and EVLA == For each observing program scheduled on an NRAO telescope, reimbursement may be requested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals. To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/. To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC. == Student Assistance for Data Reduction Visits to the DSOC == Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested. == Computing at the DSOC == A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux. For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer. Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also [[#Reservations for the EVLA site and/or DSOC|Reservations for the EVLA site and/or DSOC]]). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC. For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/. == Reservations for the EVLA site and/or DSOC == Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil. Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing requirements and the level of staff assistance needed must be specified through the online form. First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator. == Staying in Socorro == Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. == Help for Visitors to the EVLA and DSOC == We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions. == On-Line Information about the NRAO and the EVLA == NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff. = Publication Guidelines = == Acknowledgement to NRAO == Any papers using observational material taken with NRAO instruments (EVLA or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF: :''The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.'' == Dissertations == Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue. == Preprints == NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu). == Reprints == Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers. == Page Charge Support == The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary: :• When requested, NRAO will pay the larger of the following: ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s). ::– 100% of the page charge share for NRAO staff members. :• Page charge support is provided for publication of color figures. :• To receive page charge support, authors must comply with all of the following requirements: ::– Include the NRAO footnote in the text (see [[#Acknowledgement to NRAO|Acknowledgement to NRAO]]). ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254. = Documentation = Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see [[#On-Line Information about the NRAO and the EVLA|On-Line Information about the NRAO and the EVLA]]). Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the EVLA user, not listed here, are available from our website. :1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: “Synthesis Imaging in Radio Astronomy.” This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. :2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series. :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software. :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spectral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA. :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imaging under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html. :6. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS. :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/. :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983. :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. See http://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA. :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/. :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml. :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml. :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the package for data reduction is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf. = Key Personnel = {| border="1" align="center" |+ '''Table 14: Key EVLA Personnel''' !Name !Phone !Room !Notes |- | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office |- | John Benson || 7399 || 366 || Data archive |- | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms |- | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration |- | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets |- | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies |- | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software |- | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling |- | Vivek Dhawan || 7378 || 310 || EVLA commissioning |- | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director |- | Dale Frail || 7338 || 332 || Pulsars; transient sources |- | Miller Goss || 7267 || 332 || Spectral line |- | Eric Greisen || 7236 || 318 || AIPS |- | Helpdesk || 7213 || 262 || Computer Helpdesk |- | Leonid Kogan || 7383 || 312 || AIPS |- | Mark McKinnon || 7273 || 326 || EVLA Project Manager |- | Joseph McMullin || 7315 || 330 || EVLA Commissioning |- | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation |- | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA |- | George Moellenbrock || 7406 || 368 || CASA; polarimetry |- | Emmanuel Momjian || 7452 || 301 || EVLA commissioning |- | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry |- | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning |- | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging |- | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations |- | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry |- | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu |- | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu |- | James Robnett || 7226 || 258 || Computing Infrastructure Division Head |- | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line |- | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives |- | Ken Sowinski || 7299 || 375 || On-line systems |- | Meri Stanley || 7238 || 204 || Lead data analyst |- | Gustaaf van Moorsel || 7396 || 348 || User Support |- | VLA Operator || 7180 || VLA || On-duty EVLA Operator |- | VLBA Operator || 7251 || 269 || On-duty VLBA Operator |- | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling |- | Wes Young || 7337 || 378 || CASA system adminstration |} :Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is truncated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. For questions about telescope time allocation, please e-mail “schedsoc@nrao.edu”. The listed four-digit numbers are sufficient for calls made from within the DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835. = Acknowledgements = This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide. 97952c7744879a7b5a9aaf3360a3b2300c93635e 10 2 1442 1225 2013-07-18T14:42:44Z Admin 1 Changes per Gustaaf wikitext text/x-wiki {|width=100% style="background:transparent; font-size:8pt" cellspacing=10 |- |valign=top|[[Image:book.gif]] EVLA Post-Processing<br> CASA: [http://casaguides.nrao.edu '''CASA Reduction Guides'''] <br> AIPS: [[Key to Calcodes]] · [[:Category:Post-Processing|Special Considerations for EVLA Data Calibration and Imaging in AIPS]]<br> |} 021f24f956f9094763f926568bd1079e0c8773a0 14 86 1443 1433 2014-09-30T13:50:11Z Akepley 16 /* Introduction */ wikitext text/x-wiki ==Introduction== '''Note (1/25/2012): This guide has been updated for the lastest version of the OPT. Please contact Amanda Kepley (akepley@nrao.edu) for comments/questions suggestions about this guide and the NRAO helpdesk for any questions about the OPT.''' ""Note (9/30/2014): This guide no longer reflects the current state of the OPT due to limited support. Please consult the official OPT manual." The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line. The imaginary line can be close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, you need scans to set-up the attenuators and the gain slopes and scans to set the requantizer gains. The set-up scans for the attenuators need to occur only once at the beginning of a scheduling block. The requantizer gains scans need to occur after every 8-bit to 3-bit transition and every band change in 3-bit. See [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up/3bit Section 7.1.2] of the Guide to Observing with the VLA document for more details and some examples. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. A table to the right of the plot shows the rise/set times for different elevations for your source. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). You can also ask to avoid sunrise and sunset, which is important for high frequency observations. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the VLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the VLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the VLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 14dc7e470c452f5089e87900e4f6678b2bf0db9d 1444 1443 2014-09-30T13:50:25Z Akepley 16 /* Introduction */ wikitext text/x-wiki ==Introduction== '''Note (1/25/2012): This guide has been updated for the lastest version of the OPT. Please contact Amanda Kepley (akepley@nrao.edu) for comments/questions suggestions about this guide and the NRAO helpdesk for any questions about the OPT.''' ""Note (9/30/2014): This guide no longer reflects the current state of the OPT due to limited support. Please consult the official OPT manual."" The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line. The imaginary line can be close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, you need scans to set-up the attenuators and the gain slopes and scans to set the requantizer gains. The set-up scans for the attenuators need to occur only once at the beginning of a scheduling block. The requantizer gains scans need to occur after every 8-bit to 3-bit transition and every band change in 3-bit. See [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up/3bit Section 7.1.2] of the Guide to Observing with the VLA document for more details and some examples. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. A table to the right of the plot shows the rise/set times for different elevations for your source. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). You can also ask to avoid sunrise and sunset, which is important for high frequency observations. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the VLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the VLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the VLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 4491f20aa89292ef124544f8f5dcb73e82288647 1445 1444 2014-09-30T13:50:49Z Akepley 16 /* Introduction */ wikitext text/x-wiki ==Introduction== '''Note (1/25/2012): This guide has been updated for the lastest version of the OPT. Please contact Amanda Kepley (akepley@nrao.edu) for comments/questions suggestions about this guide and the NRAO helpdesk for any questions about the OPT.''' '''Note (9/30/2014): This guide no longer reflects the current state of the OPT due to limited support. Please consult the official OPT manual.''' The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line. The imaginary line can be close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, you need scans to set-up the attenuators and the gain slopes and scans to set the requantizer gains. The set-up scans for the attenuators need to occur only once at the beginning of a scheduling block. The requantizer gains scans need to occur after every 8-bit to 3-bit transition and every band change in 3-bit. See [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up/3bit Section 7.1.2] of the Guide to Observing with the VLA document for more details and some examples. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. A table to the right of the plot shows the rise/set times for different elevations for your source. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). You can also ask to avoid sunrise and sunset, which is important for high frequency observations. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the VLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the VLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the VLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 80cff6625b45eae21131b91dcf1e4ce85005805f 1446 1445 2015-09-24T14:56:07Z Emomjian 14 wikitext text/x-wiki ==Introduction== '''Note (1/25/2012): This guide has been updated for the lastest version of the OPT. Please contact Amanda Kepley (akepley@nrao.edu) for comments/questions suggestions about this guide and the NRAO helpdesk for any questions about the OPT.''' == '''Note (9/30/2014): This guide no longer reflects the current state of the OPT due to limited support. Please consult the official OPT manual.''' == The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line. The imaginary line can be close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, you need scans to set-up the attenuators and the gain slopes and scans to set the requantizer gains. The set-up scans for the attenuators need to occur only once at the beginning of a scheduling block. The requantizer gains scans need to occur after every 8-bit to 3-bit transition and every band change in 3-bit. See [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up/3bit Section 7.1.2] of the Guide to Observing with the VLA document for more details and some examples. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. A table to the right of the plot shows the rise/set times for different elevations for your source. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). You can also ask to avoid sunrise and sunset, which is important for high frequency observations. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the VLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the VLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the VLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) f36fbf8c4d1bc055a0b4a0528d5b011e742f9c4b 1447 1446 2015-09-24T14:58:09Z Emomjian 14 wikitext text/x-wiki ==Introduction== '''Note (1/25/2012): This guide has been updated for the lastest version of the OPT. Please contact Amanda Kepley (akepley@nrao.edu) for comments/questions suggestions about this guide and the NRAO helpdesk for any questions about the OPT.''' <span style="font-size:140%;> '''Note (9/30/2014): This guide no longer reflects the current state of the OPT due to limited support. Please consult the official OPT manual.''' </span> The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line. The imaginary line can be close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, you need scans to set-up the attenuators and the gain slopes and scans to set the requantizer gains. The set-up scans for the attenuators need to occur only once at the beginning of a scheduling block. The requantizer gains scans need to occur after every 8-bit to 3-bit transition and every band change in 3-bit. See [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up/3bit Section 7.1.2] of the Guide to Observing with the VLA document for more details and some examples. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. A table to the right of the plot shows the rise/set times for different elevations for your source. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). You can also ask to avoid sunrise and sunset, which is important for high frequency observations. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the VLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the VLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the VLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 3f7a86e9ad3368c0b443f6e44bd5eff7d0893070 1448 1447 2015-09-24T14:58:51Z Emomjian 14 wikitext text/x-wiki ==Introduction== '''Note (1/25/2012): This guide has been updated for the lastest version of the OPT. Please contact Amanda Kepley (akepley@nrao.edu) for comments/questions suggestions about this guide and the NRAO helpdesk for any questions about the OPT.''' <span style="font-size:140%;"> '''Note (9/30/2014): This guide no longer reflects the current state of the OPT due to limited support. Please consult the official OPT manual.''' </span> The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line. The imaginary line can be close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, you need scans to set-up the attenuators and the gain slopes and scans to set the requantizer gains. The set-up scans for the attenuators need to occur only once at the beginning of a scheduling block. The requantizer gains scans need to occur after every 8-bit to 3-bit transition and every band change in 3-bit. See [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up/3bit Section 7.1.2] of the Guide to Observing with the VLA document for more details and some examples. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. A table to the right of the plot shows the rise/set times for different elevations for your source. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). You can also ask to avoid sunrise and sunset, which is important for high frequency observations. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the VLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the VLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the VLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 4632267fa177e58caa93ffb46e7f8dd603887a39 1449 1448 2015-09-24T14:59:28Z Emomjian 14 wikitext text/x-wiki ==Introduction== '''Note (1/25/2012): This guide has been updated for the lastest version of the OPT. Please contact Amanda Kepley (akepley@nrao.edu) for comments/questions suggestions about this guide and the NRAO helpdesk for any questions about the OPT.''' <span style="font-size:140%; color=red"> '''Note (9/30/2014): This guide no longer reflects the current state of the OPT due to limited support. Please consult the official OPT manual.''' </span> The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line. The imaginary line can be close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, you need scans to set-up the attenuators and the gain slopes and scans to set the requantizer gains. The set-up scans for the attenuators need to occur only once at the beginning of a scheduling block. The requantizer gains scans need to occur after every 8-bit to 3-bit transition and every band change in 3-bit. See [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up/3bit Section 7.1.2] of the Guide to Observing with the VLA document for more details and some examples. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. A table to the right of the plot shows the rise/set times for different elevations for your source. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). You can also ask to avoid sunrise and sunset, which is important for high frequency observations. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the VLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the VLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the VLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 4e65b6e8a34c7bedf91f4f32adc84c2bb412aca0 1450 1449 2015-09-24T15:00:18Z Emomjian 14 wikitext text/x-wiki ==Introduction== '''Note (1/25/2012): This guide has been updated for the lastest version of the OPT. Please contact Amanda Kepley (akepley@nrao.edu) for comments/questions suggestions about this guide and the NRAO helpdesk for any questions about the OPT.''' <span style="font-size:140%; color:red"> '''Note (9/30/2014): This guide no longer reflects the current state of the OPT due to limited support. Please consult the official OPT manual.''' </span> The goal of this guide is to provide a brief, practical introduction to using the Observing Preparation Tool (OPT) to create scheduling blocks for the VLA. After reading this guide, you should be able to successfully create a simple VLA scheduling block. Please send any questions, comments, or suggestions about this documentation to Amanda Kepley (akepley@nrao.edu). For questions about the OPT, please contact the NRAO Help Desk via the my.nrao.edu portal. A scheduling block defines a complete set of VLA observations including your source and all necessary calibration (flux, complex gain, bandpass, etc). If you have experience with the old VLA, a scheduling block is the equivalent of an VLA "observe file" created using JObserve. Your total VLA time allocation may consist of multiple scheduling blocks; these blocks do not have to correspond to the lengths of the blocks in your disposition letter. Most scheduling blocks are scheduled dynamically based on the conditions at the VLA. This guide assumes that you understand what data is necessary to successfully calibrate your VLA data. If you need a refresher, some helpful resources are: * [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] especially [http://adsabs.harvard.edu/abs/1999ASPC..180...79F Chapter 5 -- Calibration and Editing] and [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11 -- Spectral Line Observing I: Introduction], * [https://science.nrao.edu/facilities/vla/oss Observational Status Summary], which gives an overview of the VLA including calibration and current restrictions, * [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA], which gives the details of how to observed with the VLA, and * the introduction to the [http://www.vla.nrao.edu/astro/calib/manual/ VLA calibrator manual], which gives the best practices for calibrating VLA data This quickstart guide does not include the all current observing restrictions for the VLA. These restrictions are given in [https://science.nrao.edu/facilities/vla/oss VLA observational status summary] and [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/req Chapter 5] of [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] and change frequently. Make sure you that you read and understand the current restrictions. The [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual] gives more detailed information about the OPT. Since the OPT is being actively developed, some images may not be identical to those shown here. There may also be some additional features. However, the overall philosophy and workflow in the OPT should not change. If you need more information about any of the features described in this quick start guide, consult the full [https://science.nrao.edu/facilities/vla/docs/manuals/opt OPT manual]. If you have questions about the OPT, you can submit a Help Desk ticket via my.nrao.edu. ==Overview of the Scheduling Block Creation Process== There are three steps to creating a scheduling block. # Create a catalog of the sources you would like to observe. This should include your science source(s) and your calibration sources (flux calibrator, bandpass calibrator, complex gain calibrator, etc). # Configure the VLA receivers and correlator for your proposed observations. An individual VLA receiver and correlator setup is referred to as an instrument configuration. # Create a scheduling block using the source catalog created in step 1 and the instrument configuration(s) created in step 2. While completing these steps, keep a copy of your original proposal and your disposition email handy. These documents should have all the information necessary to create your scheduling blocks. ==OPT access and layout== To access the OPT, * Login to my.nrao.edu. Your login should have already been created as part of the proposal application process. If you have forgotten your password, you can use the "Forgot password?" link to obtain your login information. * Click on the "Obs Prep" tab. * Click on the "Login to the Observation Preparation Tool" link. You should now have access to the OPT. The layout of the OPT is shown below. [[Image:opt_overview_figure.png]] The different portions of the OPT are labeled in the above image. The Menu and Navigation strips occupy the top two lines of the screen. The Menu strip contains the usual "File", "Edit", and "Help" categories. The "View" menu item allows you to select which projects you would like to be visible in the Observational Preparation part of the tool, which is handy if you've been lucky enough to have many accepted proposals. Below it, the Navigation strip allows you to navigate to the different sections of the OPT: * the Observation Preparation section where you prepare your scheduling blocks, * the Sources section where you put together your source lists, and * the Instrument Configuration section where you set up the receiver(s) and correlator configuration(s). Any system-wide messages appear in red below the navigation strip. These messages alert users to system downtime, critical bugs, etc. The item listing is in the left hand column. This list consists of either a list of scheduling blocks, instrument configurations, or source catalogs, depending on which section of the OPT you are in. The button bar above the item listing allow you to modify the item listing (e.g., cut, paste, copy, etc). You can hover the mouse over the different buttons to determine their function. Many of the button bar functions can also be accessed via the "File" and "Edit" menus. The right hand column is your main editing window. In this window, you will set up scans in your scheduling block, add source information, and edit correlator configurations. Error messages will appear as necessary in the bottom portion of your screen. The OPT automatically saves your work. You can create your scheduling block in several shorter sessions rather than having to do everything at once. ''WARNING: The OPT allows any co-investigator on your proposal to edit and submit scheduling blocks. However, only one individual can edit scheduling blocks at a time. To make sure that other members of your team are able to edit the scheduling blocks for your project, always click on the "EXIT" button in the upper right hand corner of the OPT to leave the OPT rather than just closing your browser window.'' ==Creating a Source List== To create a source list, click on "Sources" in the Navigation strip at the top of the screen. The item listing on the left hand side of the OPT will show your source catalogs. An example of this display is shown below. Your list of catalogs will be different from those shown. The only catalog that will be the same will the VLA catalog. [[Image:source_overview.png]] Each catalog contains one or more groups of sources. Groups allow you to organize your source list and make it easier to generate scheduling blocks. For example, you can include the calibrators for your science source in the same group as your science source. You can also manually add catalogs, groups, and sources using File-&gt;Create New. A newly created source only has fields for you to enter the source name and position. Your source names need to be less than 14 characters for your data to be compatible with AIPS. In addition to science targets, most observations will need three different types of calibration sources: complex gain (i.e., phase) calibrators, flux calibrators, and bandpass calibrators. We will go through the process of selecting these different types of calibrators in the OPT. This will demonstrate both how to select these types of calibrators and the different ways to select sources from catalogs using the OPT. ''Warning: Your observations may have different calibration needs than the above. It is your responsibility to make sure that your data will be adequately calibrated.'' ===Adding Your Science Targets=== The list of sources from your accepted proposal are automatically imported into a catalog. The name of the catalog is either the legacy ID for your proposal (e.g., AK774) or the new id (e.g., 12B-292). The IDs for your proposal is given on the cover page of your proposal, in the list of submitted proposals in the Proposal Submission Tool at my.nrao.edu, and in your disposition email. You can look at the parameters for or edit the sources in your catalog by clicking on the edit button to the left of the source name in the catalog as shown in the image below. It's a good idea to check the source names, positions, and velocities that have been imported from the PST. You can add a source velocity if your source doesn't have one already by clicking "Add" in the "Source Velocity" section. ''Note: You may have been awarded time for only some of your proposed sources. Regardless of the awarded time, all your proposed sources are imported in a catalog, not just those you were awarded time to observe.'' If your sources haven't been automatically imported, use the "External Search" box to find your source in NED or SIMBAD. Once you have found the appropriate source, follow the procedure outlined in the next section as to copy and paste it into your catalog. [[Image:source_edit_button.png]]<br /> You can also enter your source manually by selecting your catalog, choosing File-&gt;Create New-&gt;Source. You will need to fill in the source name, source position, and possibly the source velocities sections of the source editing screen. ===Selecting a Complex Gain Calibrator=== Complex gain calibrators are used to calibrate amplitude and phase as a function of time. They are typically point sources located close (&lt; 10 degrees) to the science source. The easiest way to select a complex gain (i.e., phase) calibrator is to click on the group containing your science source in the catalog for your observing program. To the right of your science source name in the main editing window, there will be a link to a map of the sky around your object (indicated below by an orange circle in the image below). Click on the link. [[Image:skymap_option.png]] When you click on the "Sky Map" link, a separate window similar to the following will appear below. The left hand side shows all the calibration sources near your selected source. If you hover your mouse above a particular calibration source, information from the VLA calibrator manual about that source is shown. See the VLA calibrator manual for a [http://www.vla.nrao.edu/astro/calib/manual/key.html key to the catalog listing]. Both the [http://www.vla.nrao.edu/astro/calib/manual/hints.html VLA calibrator manual] and the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] have tips for selecting complex gain calibrators. Once you're selected a suitable source, write down the name of the source. ''Warning: Clicking on the sky map will recenter it at those coordinates. Don't click on the map if you don't want to recenter it.'' [[Image:skymap_window.png]] All VLA calibrators shown on the map are listed in the VLA catalog. To add the selected complex gain calibrator to the catalog for your project, search the VLA catalog for that calibrator and then copy and paste the complex gain calibrator information. To do this, * Select the VLA catalog. This catalog should be at the top of your list of catalogs. * <p>Type the calibrator name in the search box at the top of the item manipulation region.</p> <p>[[Image:phase_calibrator_search.png]]</p> * Click search. Your desired calibrator should appear in the main editing window. * Click the check box next to the calibrator. * Select Edit-&gt;Copy-&gt;Sources. You can also use the copy button ([[Image:copy_icon.png]]) in the button bar. [[Image:phase_calibrator_copy.png]]<br /> * Select the source catalog for your program. * Select Edit-&gt;Paste-&gt;Sources. You can also use the paste button ([[Image:paste_icon.png]]) in the button bar. [[Image:phase_calibrator_paste.png]] * <p>If you have set up source groups, then when you paste your source, the OPT will prompt you to choose the correct group for your calibrator. Select the correct group and click "Add To Group(s)" button. Grouping sources makes the book-keeping easier when creating scheduling blocks. If you want, you can also skip this step by clicking the "Skip This Step" button.</p> <p>[[Image:phase_calibrator_select_group.png]]</p> ===Selecting a Flux Calibrator=== All the flux calibrators for the EVLA are located in the VLA catalog in the group "VLA Flux Cal". If you hover the mouse in the Flux/Structure column for a particular source, you can get the VLA calibrator manual information for that calibrator. <p>[[File:Flux_cal.png.png]]</p> Choose an appropriate flux calibrator for your observations. Then use the copy and paste method detailed in the "Selecting a Complex Gain Calibrator" section to add your flux calibrator to your project catalog. ===Selecting a Bandpass Calibrator=== Objects like bandpass calibrators, which need to be very bright and up at the same time as your program source, require a more complex catalog search. (Sometimes your flux calibrator can be a suitable bandpass calibrator.) To do this type of search, * Click on the "Advanced Search" link in the left hand column. * Select the catalog you want to search. Here, we select the VLA catalog to search the VLA calibrator manual. * Select the parameters you want to search on. * Click on the "Search" button at the bottom of the page. When you've found an appropriate source, you can copy and paste it into your program catalog using the same procedure as for your complex gain calibrator. In the example below, the VLA catalog is being searched for sources brighter than 5.0 Jy in C band with RA's between 0 and 6h and declinations greater than 30 degrees. <br />[[Image:complex_catalog_search.png]]<br /> ==Setting up the Instrument Configuration== Configuring the receivers and correlator properly is one of the most important and difficult aspects of creating a scheduling block. Spectral line users should take extra care to ensure that their correlator set up is correct! ''Incorrect correlator setups may make doing your proposed science impossible.'' If you have any questions, please consult the NRAO helpdesk via the my.nrao.edu portal. The VLA WIDAR correlator is extraordinarily flexible. This quickstart guide only discusses the general correlator modes. Although the interface for configuring the correlator is similar for the shared risk modes, shared risk users should consult VLA staff to create their correlator set up. Correctly configuring the correlator requires a basic understanding of how the correlator works. There are two modes for the correlator: 3-bit and 8-bit mode. In 8-bit mode, you get two 1-GHz wide baseband pairs. In 3-bit mode, you get four 2-GHz wide baseband pairs. For general observations, you can use 16 sub-bands per baseband pair. The individual sub-band resources can be stacked to provide higher spectral resolution. You can also change the number of channels in an individual sub-band by changing how many polarization products are recorded. If you are not interested in polarization, then dual polarization is a good setting. For both 8-bit and 3-bit observations, the input frequencies for each baseband is divided up into 128MHz chunks that are fed into the correlator. ''A critical point to remember is that 128 MHz boundaries (AKA 128 MHz suckouts) have reduced sensitivity. Do NOT place lines of interest on the 128 MHz boundaries.'' See [https://science.nrao.edu/facilities/vla/docs/manuals/opt/using-the-rct/test/example-of-a-resource-catalog-the-nrao-defaults-catalog the OPT manual] for a more complete description of the correlator modes and Chapter 15 of the [https://science.nrao.edu/facilities/vla/oss Observational Status Summary]. ===Configuring for "Continuum" Observations=== If you observing continuum, there are default continuum configurations available for you to use in the "NRAO Defaults" group. These setups will also work for spectral line work provided that they adequately resolve your line and you don't want to use Doppler setting (see next section). The default continuum configuration group includes both 3-bit and 8-bit configurations. The catalog name is currently "D/Any config 13A". Use the catalog comments to determine which one you want to use. You can use the copy and paste procedure to add one of these configurations to your personal catalog. ''Currently copying a individual configuration out of a resource group doesn't work. However, you can still copy the setting if you select the entire catalog and then the individual resource.'' Once you copy and paste the configuration into your personal catalog, you can change the configuration name, the default baseband frequencies and polarization. You can change the name of your catalog in the basics tab. Changing it to a unique descriptive name will help with creating and checking scheduling blocks. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. [[Image:initial_continuum_changes.png]] To change your baseband frequency, change the frequency for the center of the baseband in the basebands tab. Do not select Doppler Setting! The plot above should change to show your current baseband tuning. The 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the basebands. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed. [[Image:baseband_continuum_changes.png]] To change the polarization of the sub-bands, go the subbands tab. From there, select all the sub-bands, choose bulk-edit sub-bands, select the desired polarization, and click okay. You need to repeat this for each baseband. [[Image:polarization_continuum_changes.png]] If you're doing spectral line work using the standard "continuum" setups, enter the sky frequency plus an offset equal to half the sub-band width into the "SB Center Freq" boxes. IT IS VERY IMPORTANT THAT YOU DO NOT ENTER THE SKY FREQUENCY INTO THE "SB Center Freq" BOX. IF YOU DO THIS, YOU WILL PLACE YOUR LINE AT THE BOUNDARY OF TWO 128MHz FILTERS WHERE THE SENSITIVITY IS SIGNIFICANTLY DECREASED AND CALIBRATION BECOMES MUCH MORE DIFFICULT. Adding an offset of 1/2 the width of an individual sub-band places your line in the center of a sub-band. You can calculate the sky frequency manually using the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool. Finally, if you're made any changes validate your correlator setup in the validation tab. This tab gives an overview of the correlator resources you're using and the setup of the correlator. [[Image:validation_continuum_changes.png]] ===Configuring for Spectral Line Observations=== If you are observing spectral lines, then you need to account for the motion of the Earth and the motion of the source when determining where to place lines (see [http://adsabs.harvard.edu/abs/1999ASPC..180..201W Chapter 11] of [http://adsabs.harvard.edu/abs/1999ASPC..180.....T Taylor, Carilli, and Perley. 1999, Synthesis Imaging in Radio Astronomy II] for a review of the different velocity systems). The old VLA used "doppler tracking". In doppler tracking mode, the sky frequency was changed for every scan so that the same rest frequency would always be in the same channel. The current VLA does not doppler track. Instead, you can either enter the sky frequency manually or use the doppler setting mode. Choosing whether to enter the sky frequency manually or to use doppler setting depends on your science and desired correlator settings. If your sub-bands are much wider than 30 km/s, then correction for the motion of the Earth is much smaller than the width of your sub-bands. Therefore, using doppler setting does not get you much and it is best to enter the sky frequency manually. You can account for the change in the sky frequency of the line with time in post-processing using a task like ''cvel'' in AIPS or CASA. Where doppler setting comes in handy is with narrower sub-bands (less than a few x 30 km/s). These sub-bands are small enough that the motion of the Earth may shift your line out of the sub-band. Doppler setting allows you to submit one scheduling block and have the line automatically be placed in the right spot. Otherwise, you would have to recalculate the sky frequency manually every few weeks and resubmit your scheduling block to account for the motion of the Earth. Note that doppler setting is more difficult to set up in the correlator (most of this complexity is hidden from the end user). To enter the sky frequency manually, you can use the [http://www.vla.nrao.edu/astro/guides/dopset/ dopset] tool (http://www.vla.nrao.edu/astro/guides/dopset/) to calculate the sky frequency of your line. You need to enter the source position, the source velocity, the rest frequency of the line, and the observation date and time. The webpage then returns the appropriate sky frequency to enter into the instrument configuration editor. To get the LST day numbers, you can consult the [http://www.vla.nrao.edu/cgi-bin/schedules.cgi monthly EVLA schedule]. The LST days are given to the right of the schedule grid. To get the LST day numbers more than a month in advance, you can extrapolate from this schedule. Doppler setting mode is configured using the instrument configuration editor. In doppler setting mode, the appropriate sky frequency is automatically calculated once at the beginning of the scheduling block. However, unlike doppler tracking mode, the sky frequency is the same throughout the entire observation; it does not change with every scan. The process for configuring the WIDAR correlator for spectral line observations is iterative. The basic idea is to get the receiver and spectral windows containing lines set up, then go back and set the frequencies for the centers of each baseband, doppler setting (if desired), and any continuum windows. To set up your spectral line configuration: * Create a new spectral line resource catalog by selecting File-&gt;Create New-&gt;Catalog. * Create a new spectral line resource by selecting File-&gt;Create New-&gt;8-bit Instrument configuration. General spectral line observations that need more than the standard 3-bit continuum configuration are currently required to use 8-bit mode. A new instrument configuration window should pop up. The window has three parts: 1) a plot showing the current correlator configuration, 2) a table showing the correlator resources used, and 3) a series of tabs for configuring the correlator. In the plot, the 3 dB edges of the band are indicated by dotted lines. The green lines show the full sensitivity range of the band. If you hover your mouse over a particular baseband, the center frequency of that baseband will be displayed [[Image:spectral_line_correlator_window_overview.png]] * The first tab you should see is the basics tab. Here you can name your configuration, select the receiver, and desired integration time. Please use the default integration time for your observing frequency and configuration. Do not use integration times shorter than the default unless it is necessary for your science. Shorter integration times lead to higher data rates, which is given in the "Correlator Setup" section under the integration time. <p>[[Image:spectral_line_correlator_basics.png]]</p> * Next go to the lines tab. In this tab, enter the information on the source and spectral lines you would line to observe. You can import source information from the catalog you created earlier to avoid having to retype things and introduce new errors. When you have multiple lines, copying the previous spectral line prevents you from having to re-enter defaults like the velocity definition and source velocity. The spectral line information is used generate sub-bands of the appropriate bandwidth and channel resolution to observe that line. You can change the bandwidth and channel resolution for the individual line sub-bands later on in the sub-band configuration tab. The lines that you've added will show up in the overview plot above the tabs.<p> [[Image:spectral_line_correlator_line_setup.png]]</p> * Now on to the basebands tab! In this tab, center the basebands approximately where you want them, but '''do not select Doppler Setting yet.''' It's easier to set up the sub-bands before Doppler Setting has been activated. The base bands should roughly overlap your lines. <p>[[Image:spectral_line_correlator_baseband_setup.png]]</p> * It's time for the fun part -- generating your spectral windows. To do this, go to the line placement tab, select the desired line, and click generate. The pop-up dialog box will ask you which baseband you want to add it to. If your line is already in a particular baseband, it should select that baseband automatically. Select the appropriate baseband (if not already selected) and click Okay. If you click generate multiple times you will get multiple windows. You can just delete the extra windows in the sub-band configuration window. As they are generated, the sub-bands will show up in the overview plot. <p>[[Image:spectral_line_correlator_generate_spws.png]]</p> * The sub-bands you generated will show up in the subband tab. Make sure all your line sub-bands have the appropriate properties. The properties of the generated sub-band (including the central frequency) can be changed in this time if they are not appropriate for your science. If you want to change sub-band properties, you need to do it separately for each baseband, i.e., you can't change sub-band properties across all base bands. Remember that you can't place sub-bands across the 128 MHz filter boundaries! If you would like to add continuum sub-bands, hold off for now. It's easier to do this after you're set up doppler setting.<p>[[Image:spectral_line_correlator_subband_window.png]]</p> * If you're getting errors about your sub-bands being too close to baseband boundaries or if you don't like where they are placed in the baseband, go back to the baseband tab and change the center of the baseband to a more appropriate value. This process may take some fiddling. The overview plot and the plot in the sub-band tab can help guide your sub-band placement. * Once you're happy with the placement of the line sub-bands, you can turn on doppler setting if desired. Technically, the correlator can only do Doppler setting for individual basebands, not individual sub-bands. Therefore, you need to use a line to Doppler set on for each baseband. This line can either be one of your science lines or an imaginary line. The imaginary line can be close to the center of your baseband or the average of your line frequencies depending on what makes the most sense for your data. Consult the NRAO help desk if you have any questions about the best way to Doppler set for your program. <p>[[Image:spectral_line_correlator_doppler.png]] </p> * Finally, you can go back and add your continuum sub-bands in the sub-bands window. While the narrower sub-bands can be flexibly tuned within the 128 MHz boundaries, the larger spectral windows (e.g., 64 MHz and 128 MHz) have typically fixed placements. You can automatically fill the rest of the baseband with continuum windows by selecting "Fill subbands". To edit the properties of all the sub-bands for each baseband, select the continuum sub-bands, and then select "Bulk Edit Selected Subbands". Remember that you need to do this for each baseband. <p>[[Image:spectral_line_correlator_continuum_windows.png]]</p> * Once you're happy with your setup, validate it in the validate tab. This tab gives lots of information about the correlator including your data rate, the correlator resources used, where your sub-bands are placed in the baseband, and which baseline boards this configuration will use. It's a good idea to use the data rate to estimate the total data volume for an individual scheduling block and for your entire project. The current VLA produces much more data than the old VLA. Large data files take a long time to download and process. For example, a 32GB data file takes about 8 hours to download over high-level broadband (10Mb/s) and about 1hr of computing time to run through a calibration script using CASA 3.4 on an computer with 12 processors, 24GB of RAM, and 2TB of disk space. Therefore, it is important to minimize your file sizes as much as possible. If you're not a shared risk observer, do not exceed the data rate limit set in the Observational Status Summary. Also make sure to resolve any errors that pop up at the bottom of the screen associated with your configuration. Do not submit a scheduling block with an invalid correlator configuration. <p>[[Image:spectral_line_correlator_validation.png]]</p> ===Configurations for Pointing Scans=== At high frequencies, you need to periodically obtain pointing solutions. These pointing solutions are typically obtained at C or X band. There is a NRAO-provided catalog of the default pointing setups. To access this catalog, * Go to the "Instrument Configuration" section of the OPT. * Click on the "NRAO Defaults" catalog on the left hand side of the screen. * <p>Click on the "Pointing setups" group in the left hand side of the screen.</p><p> [[Image:pointing_setups.png]]</p> To aid in creating your scheduling block, it's useful to copy and paste the desired pointing setup into your personal instrument configuration catalog using the copy and paste method outlined in the calibrator selection section. ==Putting It All Together: Creating a Scheduling Block== Before creating your scheduling block, please see the [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide Guide to Observing with the VLA] for the latest VLA observing restrictions. These restrictions change frequently, so check back often to make sure you have the latest information. The most basic restriction on your scheduling block is its length. The length of your scheduling block needs to be a multiple of 15min. Your scheduling blocks do not need to be the same length as the blocks listed in the email from the schedsoc committee. They can be as long or as short as necessary. Longer blocks (&gt;4 hrs) may have significant changes in observing conditions, while shorter blocks (&lt;2 hrs) require higher overhead to obtain adequate calibration observations. To choose an appropriate scheduling block length, you need to balance your project's scientific needs and its science rating with your allocated VLA time as a function of LST. Projects with lower ratings in popular LST ranges may want to schedule in shorter blocks to increase the probability that they are observed. The [https://science.nrao.edu/facilities/evla/schedules/schedsoc schedsoc] page includes a link to the VLA pressure histogram for your observing semester. In 8-bit mode, the VLA requires an initial dummy scan to set the attenuation levels. The dummy scan should have a length of at least 1 minute and should have the same correlator set up as the first "real" scan. At high frequencies where pointing is required (&gt;18 GHz), two dummy scans are required. Each should have a length of at least 1min. The first has the same correlator set up as the first "real" scan and the second has the same correlator set up as the pointing scans. If you have more than one science correlator setup, you need additional dummy scans for each correlator set up. In 3-bit mode, you need scans to set-up the attenuators and the gain slopes and scans to set the requantizer gains. The set-up scans for the attenuators need to occur only once at the beginning of a scheduling block. The requantizer gains scans need to occur after every 8-bit to 3-bit transition and every band change in 3-bit. See [https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up/3bit Section 7.1.2] of the Guide to Observing with the VLA document for more details and some examples. To create a scheduling block, * Click on the "Observation Preparation" link in the navigation strip at the top of the screen. * In the item listing on the left hand side of the OPT, click on the "+" side to the left of the project code for your proposal. Project codes are designated with an light brown box with a P inside it and have a format like 12A-186, ([[Image:project_code.png]]). The code is listed on your proposal cover sheet and on the submitted proposals list in the Proposal Submission Tool. * In the level below the project code level, click on the plus sign to the left of project code to see the configurations for which you have be awarded time. These are referred to as program blocks and are designated by a blue box with the letters PB inside ([[Image:program_block.png]]). The configuration for the program block will be indicated by a letter to the right of the icon. If your project can be observed in any configuration, you don't need to click on the relevant configuration to start to create a scheduling block.<br /> * Finally, in the level below the program block level, click on the "+" sign next to the relevant configuration to start creating a scheduling block. Scheduling blocks are designated by a green box with the letters "SB" inside it ([[Image:scheduling_block.png]]). There should already be a scheduling block stub available. If you want to create a new scheduling block, go to File-&gt;Create New-&gt;Scheduling Block. The item listing on left hand side of the OPT display should now look similar to following: <br /> [[Image:opt_block_structure.png]] <br /> The top level is the project code level, the next level down is the program block level, and the final level is the scheduling block level. The first page you will see after you have created a new scheduling block is the 'Information' tab. There are three items you will likely change on this page: * the name of the scheduling block, * the LST start range, and * the scheduling constraints/conditions. The name of the scheduling block and the scheduling constraints/conditions boxes are self-explanatory. The LST start range is a little more subtle. To determine the LST start range for your scheduling block, * Click on the "Sources" link in the navigation section of the OPT. * Click on the catalog with your program sources in it. * <p>Click on the edit button for one of the source(s) that you will be observing in your scheduling block.</p> <p>[[Image:source_edit_button.png]]</p> * <p>In the editing window, click on the image tab to see the the elevation and azimuth curves for your source.</p> <p>[[Image:source_image_tab.png]]</p><p>[[Image:source_elevation_plot.png]]</p> * Note the range of LSTs that your source is at an appropriate elevation/azimuth for observing. A table to the right of the plot shows the rise/set times for different elevations for your source. Repeat this process for all sources in your scheduling block including calibrators. The appropriate LST start range for your scheduling block is the range of LSTs during which all your science and calibration sources are up for the entire scheduling block. Make sure to take into account the length of your scheduling block. If all your sources are up from 10h to 18h LST and your scheduling block is 4h long, then the appropriate LST start range is 10h to 14h. Now go back to your scheduling block in the "Observation Preparation" section of the OPT. In the information tab for your scheduling block, * uncheck the box for "no constraint" * <p>input the appropriate LST start ranges for your scheduling block making sure to take into account the length of your scheduling block.</p><p>[[Image:lst_start_range.png]]</p> You can input multiple LST ranges by clicking the "Add" button under the initial LST range. This is handy for avoiding LST ranges where your source is too high (elevation &gt; 80 degrees). You can also ask to avoid sunrise and sunset, which is important for high frequency observations. One more thing to note in the Information tab is the "Count" entry under the scheduling block name. If you would like a scheduling block to execute multiple times, you can change the count from 1 to your desired number of executions. When submitting a scheduling block with multiple executions, you might consider submitting a block with a count of 1, checking the data after it has run, and then copying and pasting that scheduling block to a new scheduling block with a count greater than 1. Note that you cannot change the count on a scheduling block once it has been executed on the telescope. Now that you have filled in the information tab, you are ready to add scans to your scheduling block. You can either create individual scans or you can create a number of scans at one time. To create an individual scan, click on the scheduling block and go to File-&gt;Create New-&gt; Scan. Once you have created a scan, you should see something like the following: <br /> [[Image:new_scan.png]] The values that you most likely will want to change are shown in orange. Most of the values are self-explanatory. For the Target Source and Hardware Setup, you should click "Change" and then navigate to the appropriate catalogs and select the desired source or instrument configuration. You also should change the "intents" value to reflect the purpose of the scan. This will help with pipeline processing. If you are doing reference pointing, you should select "Interferometric pointing" under Scan Mode for pointing scans and select "Apply Last?" under Reference Pointing for source scans. The "Bulk Scan Creation" allows you to create multiple scans with similar properties in one step. For example, if you have a group of sources that all have the same instrument configuration, you could use the "Bulk Scan Creation" tool to create multiple scans each with a different source but the same instrument configuration. This allows you to avoid having to create a new scan for each source, which could be very time-consuming depending on the number of sources. Another use of "Bulk Scan Creation" might be to import different mosaic pointings, but give them the same duration, scan intent, and instrument configuration. To use "Bulk Scan Creation", * click on the "Bulk Scan Creation" tab. * Fill in the desired values in the "Create Scans From Source Group" section. The most likely of these are the Hardware setup and the Source group which are circled below in orange. * Click on create scans. * <p>Edit the individual scans as needed.</p> <p>[[Image:bulk_scan_creation.png]]</p> Once you have created all the necessary scans for your scheduling block, you can arrange them using either the cut and paste buttons or the arrow buttons at in the button bar to move scans around. The plus and minus buttons allow you to expand and collapse the scheduling block structure. Scan loops are a useful way to repeat groups of observations. For example, you could put your science source and complex gain calibrator in a loop. Then you only need to create two scans and a loop rather than 2 times n scans. Unlike the old VLA, a scan loop for the VLA does not have any additional overhead. To create a scan loop, go to File-&gt;Create New-&gt;Scan Loop. You have a choice of where to put the scan loop (before or after a scan or possibly in another scan loop). If you click on the scan loop, you should see a screen similar to the following where you can set the properties of the scan loop. <br /> [[Image:scan_loop.png]] If you want to change many scans at once, the "Bulk Scan Edit" option is very handy. To bulk edit scans, go to the "Bulk Scan Edit" tab in either the top level of a scheduling block or the top level of a loop. The bulk edit will only affect the scans in the structure that you are in. For example, if you click on "Bulk Scan Edit" in a loop, you will only bulk edit scans in that loop. If you click on "Bulk Scan Edit" in the top level of your scheduling block, it will bulk edit all the scans in your scheduling block. When you select the "Bulk Scan Edit" tab, you will see something like the following: <br />[[Image:bulk_scan_edit_select.png]] To bulk edit scans, * Click on the fields that you would like to select on. In the example above, we have selected on the source name. * Click select. * <p>Now click on the fields that you could like to change. In the example below, we have decided to change the resource used for this source.</p><p> [[Image:bulk_scan_edit_change.png]]</p> * Now click "Update" * <p>A list of the changes that will be made will appear in the editing window. Click confirm to accept this changes and back to make further change to your selecting and editing parameters.</p><p>[[Image:bulk_scan_edit_confirm.png]]</p> ''Warning: If you change one of your source catalog entries or instrument configurations in their respective sections, these changes do NOT automatically propagate to your scheduling block. You should bulk edit the affected scans in the observing preparation section of the OPT to update their source information or instrument configuration. The old and new sources and/or instrument configurations may have the same name. However, giving different names to the new sources and instrument configurations will help when double-checking your scheduling block.'' ==Checking Your Scheduling Block== Congratulations on creating a scheduling block! Your scheduling block is automatically saved. If you decide to exit the OPT before submitting your scheduling block, the draft scheduling block will still be in the OPT when you return. Before submitting your scheduling block, you should double-check and triple-check it to ensure that it is correct. Since your script will be run by the operator, there is no chance to change things once your scheduling block is on the telescope. The Report tab on the top level of your scheduling block contains a wealth of information on your scheduling block. We list some important items to check for the scheduling blocks below. If you have any questions about your scheduling block, please contact the NRAO helpdesk via the my.nrao.edu portal. * Information tab: ** Is your LST start range appropriate? The easiest way to test this is use the "Assumed schedule start" field in the "Report" tab. You should step through the range of acceptable LST start times in half hour increments. Check that your sources are high enough and that you have enough time on source for each LST. ** Have you selected the appropriate scheduling constraints/conditions for your observing frequency? * Report tab: ** Have you followed all the VLA observing constraints listed on the [https://science.nrao.edu/facilities/evla/observing/opt OPT webpage] including dummy scans, total time on first source, appropriate first and last sources, and scan lengths? ** Is your scheduling block the correct length (multiple of 15 minutes)? You may need to enter an assumed schedule start time and click on update for the total scheduling block time to be updated. ** Is your source position correct? ** Is the position in the correct coordinate system (J2000 vs. B1950)? ** Is your source velocity correct? ** Is the source velocity in the correct coordinate system? ** Is your correlator configuration correct? *** Do you have an appropriate integration time? *** Have you selected the correct frequency? *** Have you selected "sky" vs. "rest" frequency as appropriate? *** Did you avoid placing your lines on sub-band edges? To see the frequency range and center frequency of each sub-band, select "Show all sub-bands" in the "Instrument Configuration" section. *** Is your total bandwidth correct? *** Is your channel width correct? *** Is your number of polarization products correct? ** Do you have the correct amount of time on source for all sources? ** Do your sources have the right instrument configurations associated with them? ** Are the scan intents set correctly for each source? ** At high frequencies, have you applied the reference pointing results to your source scans? ** Have you included all the appropriate calibration data (flux calibrator, complex gain calibrator, bandpass calibrator, etc)? ** Do you have enough time on the calibrator sources? ** Are your calibration sources bright enough? ** Is your complex gain calibrator approximately a point source at your observing frequency? This information is given in the VLA calibrator manual as the "calcode". See the [http://www.vla.nrao.edu/astro/calib/manual/key.html VLA calibrator manual]for more details. ** Do your polarization calibration sources have the appropriate polarization properties at your observing frequencies? ** When stepping through the list of LST start ranges, do you have enough time on source for all scans (especially pointing and calibration scans)? ** When stepping through the list of LST start ranges, do you have any antenna shadowing warnings? Antenna shadowing occurs when one antenna blocks light going toward a second antenna. It is especially common in compact configurations for low declination sources. If a significant fraction of your antennas are shadowed for a particular LST start time, you may want to consider excluding that time from the range of acceptable LST start times. Shadowed antennas can be flagged. ** When stepping through the list of LST start ranges, do your sources have elevations of between 30 and 80 degrees if observing at high frequencies?<br /> ** When stepping through the list of LST start ranges, do your sources have elevations between 9 degrees and 80 degrees if observing at low frequencies? ** Are you in danger of hitting a wrap limit? The cabling inside the antennas has a finite length. This limits the angular distance that the antenna can rotate. When an antenna has rotated the maximum angular distance, it hits a "wrap limit" and has to rotate in the other direction to unwrap the cabling. The unwrapping process is slow and can take a significant amount (~10min) of time. Since the VLA is dynamically scheduled, it is difficult to know where the antenna will be at the beginning of your scheduling block. To avoid hitting a wrap limit in the middle of your scheduling block, it is advisable to set the antenna wrap direction in the first 4-5 scans using the "antenna wrap" box in the scan information window editing window and assign some extra time to your first "real" scan. Then any unwrap will occur during the beginning of your block. The report tab will inform you whether or not the antenna is wrapping in the middle of a scheduling block. ==Submitting Your Scheduling Block== To submit a scheduling block, click on the validation and submission tab in the "Observation Preparation" section. First, you need to validate the scheduling block. Once the scheduling block has been validated then you can submit it using the submit button at the bottom of the page. You can request help with your submission using the link to the NRAO helpdesk in the middle of the page. To retract a scheduling block after it has been submitted, go to the validation and submission tab for the scheduling block and click "cancel". Once your scheduling block has been executed, you should receive an email from the VLA Operator letting you know that the data has been taken and is available in the archive. You and your collaborators can configure your email settings by clicking on the project code on the left than side of the OPT and going to the "Principle Investigator and co-authors" section. The Executions tab on the scheduling block gives you information on when the block was executed. ==About This Document== This document was originally created in Evernote, exported as HTML, and converted to mediawiki format using the Perl module HTML::WikiConverter. The illustrations were created using Skitch. ---- --[[User:Akepley|Amanda Kepley]] 13:01, 22 January 2013 (PST) 7bf90cfeb962de2f810ab38da78eb0063539efda